Presolar grains from meteorites

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INVITED REVIEW

Presolar grains from meteorites: Remnants from the earlytimes of the solar system

Katharina Lodders a,* and Sachiko Amari b a Planetary Chemistry Laboratory, Department of Earth and Planetary Sciences and McDonnell Center for the Space Sciences, Washington University, Campus Box 1169,One Brookings Drive, St. Louis, MO 63130, USA

b Department of Physics and McDonnell Center for the Space Sciences, WashingtonUniversity, Campus Box 1105, One Brookings Drive, St. Louis, MO 63130, USA

Received 5 October 2004; accepted 4 January 2005

AbstractThis review provides an introduction to presolar grains – preserved stardust

from the interstellar molecular cloud from which our solar system formed – found in primitive meteorites. We describe the search for the presolar components, the currently known presolar mineral populations, and the chemicaland isotopic characteristics of the grains and dust-forming stars to identify thegrains’ most probable stellar sources.

Keywords: Presolar grains; Interstellar dust; Asymptotic giant branch (AGB)stars; Novae; Supernovae; Nucleosynthesis; Isotopic ratios; Meteorites

1. Introduction

The history of our solar system started with the gravitational collapse of aninterstellar molecular cloud laden with gas and dust supplied from dying stars.The dust from this cloud is the topic of this review. A small fraction of this dustescaped destruction during the many processes that occurred after molecular cloud collapse about 4.55 Ga ago. We define presolar grains as stardust thatformed in stellar outflows or ejecta and remained intact throughout its journeyinto the solar system where it was preserved in meteorites.

The survival and presence of genuine stardust in meteorites was not expectedin the early years of meteorite studies. In the 1950s and 1960s, models of solar system formation assumed that the matter from the presolar molecular cloud was

processed and homogenized (e.g., Suess 1965, see also Fegley 1993). Most of this matter accreted to the Sun and less than about one percent remained to form

*Corresponding author.e-mail address: [email protected]

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the planets, their satellites, and other small objects (asteroids, Kuiper-Edgeworth- belt objects). During collapse and accretion of matter towards the cloud center,gravitational heating vaporized presolar solids, and it was generally assumed thatthis process resulted in a relatively homogeneous solar nebular gas made of evaporated presolar solids and presolar gas. Upon cooling of the solar nebula,

new condensates appeared whichaccumulated to form the solid bodiesin the solar system (Fig. 1). In thisvery simplified picture, all matter from the presolar cloud would bechemically and isotopicallyhomogenized, and no record aboutthe mineralogy of presolar solidswould remain.

Long before presolar grains werediscovered, Cameron (1973)speculated on the question“Interstellar grains in museums?” andconcluded that primitivecarbonaceous chondrites may harbor

presolar grains. Indeed, presolar grains were incorporated intometeorite parent bodies (smallasteroid-size objects). Chemical andmetamorphic processes on the leastmetamorphosed meteorite parent

bodies were apparently mild andmore or less non-destructive to thegrains, so upon meteorite deliveryfrom a parent body to Earth, the

presolar grains awaited discovery.There are several motivations to

study presolar dust. Fresh dust andgas are continuously supplied to theinterstellar medium (ISM), mainly bymass-loss from red giant stars and byexploding novae and supernovae,which is the way such stars enrich the

ISM with their nucleosynthetic products over time. Hence the presolar gr ainsfrom meteorites could provide a glimpse into the dust population thataccumulated in the presolar molecular cloud. However, the known presolar minerals are very likely a biased sample of solids from the presolar molecular

Fig. 1. The end of red giant stars andsupernovae is accompanied by production of dust grains that also entered the molecular cloud from which our solar system formed. Atiny portion of original stardust survived the

passage through the ISM and the eventsduring solar system formation. After extraction of presolar dust from meteoritesand examination by various types of instruments, we obtain information about thegrains’ parent stars

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cloud because of their different stabilities against physical and chemical processing in the ISM, in the solar nebula, on meteorite parent bodies, and in thelaboratory during grain isolation procedures.

Another reason to study presolar grains is that their physical, chemical andisotopic signatures record nucleosynthetic and grain formation processes of avariety of stars. Grain morphologies and compositions may reflect formationconditions in stellar environments, and the grains’ chemical and isotopiccompositions are firm tests of stellar evolution and nucleosynthesis models.Hence the microscopic grains reveal more details about these processes thancurrently possible by astronomical observations with even the best telescopes.

One way to find the nm- to µm-size presolar grains in meteorites is to searchfor variations in isotopic compositions because several elements in certainmeteoritic components are clearly different from the terrestrial isotopic referencecompositions. Many of these “isotope anomalies” cannot be easily explained bymass-fractionation processes or by decay of longer-lived radioactive isotopes andtherefore suggest the presence of exotic, possibly presolar, phases.

Even in the 1950s and 1960s there were indications that the solar nebula wasnot homogeneous. For example, Boato (1953, 1954) and Briggs (1963) founddeuterium enrichments in carbonaceous chondrites that were not easy to explain

by chemical mass-fractionation processes known at that time. Clayton (1963)encountered a similar problem in trying to explain large 13C enrichments incarbonate from the Orgueil meteorite and thought that incompletehomogenization of matter from different nucleosynthetic sources could beresponsible. Unusual isotopic compositions in noble gases from meteorites werediscovered for Xe (Reynolds and Turner 1964), and Ne (Black and Pepin 1969).

However, the discovery of widespread isotopic variations in O, the major rock-forming element, by Clayton et al. (1973) provided the strongest evidenceof incomplete homogenization of matter in the solar nebula. In particular, the so-called calcium aluminum-rich inclusions (CAIs) found in chondritic meteoritesshow characteristic enrichments of around 5% in 16O relative to terrestrial rocksand other major meteorite components. These enrichments cannot be derived bynormal mass-dependent fractionation processes in the solar nebula and requireanother explanation. Clayton et al. (1973) suggested that 16O, a major productfrom supernovae, was present in a reservoir in the solar nebula (see, however,alternative explanations of the 16O enrichments by mass- independentfractionation processes by Thiemens and Heidenreich 1983, Clayton 2002).

Further support that supernova and other stellar debris may have been presentin the presolar molecular cloud came from observations of isotopic variations inother abundant rock-forming elements such as Ca, Ti, and Cr in CAIs (Clayton etal. 1988). The variations in Ca and Ti isotopic compositions in CAIs are typicallyobserved for neutron-rich isotopes such as 48Ca and 50Ti, which are signaturesfrom nucleosynthesis operating in supernovae, but the size of these anomalies

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(0.01% to 0.1%) is quite small (Clayton et al. 1988). Thus, although CAIs oftencarry isotopic anomalies, many researchers do not regard them as pure presolar

phases, and the origin of isotopic anomalies in CAIs is rather vaguely assigned toa possible incorporation of some presolar components.

In contrast, bona-fide presolar grains are identified by their often huge,orders-of-magnitude-ranging excesses or deficits in isotopic compositions(relative to normal terrestrial isotopic abundances). A presolar phase may have anoutstanding isotopic anomaly in at least one element, but often isotope anomaliesare observed simultaneously in several major elements such as C, N, O, and Siwithin a single presolar grain. In essentially all cases, the observed isotopicanomalies can only be explained by nucleosynthetic processes, which means thatthese grains formed near the sites of nucleosynthesis. This leaves the stars toimprint their chemical and isotopic signatures onto the grains. The principalenvironments for grain formation are the circumstellar shells of red giant starsand asymptotic giant branch (AGB) stars, and supernova ejecta.

The major and trace element chemistry can also be used to identify presolar grains. For example, the occurrence of reduced grains such as graphite andsilicon carbide in meteorites consisting mainly of oxidized rock and hydroussilicates (e.g., CI and CM chondrites) is a rather unusual paragenesis, and is mosteasily explained by an external source of reduced grains. However, the presolar grains are rather small, and carbonaceous dust also may have formed in the solar nebula, so major element chemistry alone cannot provide a strong constraint onthe origin of the reduced dust. On the other hand, trace element abundances insome reduced grains cannot be explained by chemical fractionation duringcondensation from a gas of solar composition. Instead, trace elements in non-solar proportions are required at the grain sources, and several types of starsshow non-solar trace element abundances due to nucleosynthesis.

Here we review some of the history of the search and discovery of presolar dust, the minerals already identified, and what other minerals are to be expected.We then describe dust-producing stars and how certain grain types can be relatedto them. Since the discovery and isolation of presolar grains by Lewis et al.(1987), much work on presolar grains has followed (see reviews by e.g., Andersand Zinner 1993, Bernatowicz and Zinner 1997, Ott 1993, Zinner 1998, 2004).

2. The search for presolar grains

Historically, noble gas studies played an important role in cosmochemistry.

The noble – or rare –gases are literally rare in many meteorites but even a smalladdition of a trapped or adsorbed component whose isotopic composition isvastly different from “normal” is easily detected. Because most of the noblegases have more than two stable isotopes and an anomalous composition may not

be restricted to a single isotopic ratio, the term “component” is used to describe

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(Jungck 1982). This indicated the presence of almost pure 22 Ne in some phase of the separates and supported the idea of a direct nucleosynthetic origin fore Ne-E,as suggested by Black (1972). A possible source of “pure” 22 Ne is from decay of 22 Na (t1/2= 2.6 a), and with such a short half-life, 22 Na must have beenincorporated into a carrier near its stellar source to allow in-situ decay (Black 1972, Clayton 1975, Eberhardt et al. 1981, Jungck 1982).

2.2. Presolar signatures in Xe isotopes: Xe-HL, Xe-S

With nine stable isotopes, the Xe systematics in meteorites and planetaryatmospheres is quite complex so that Reynolds (1963) coined Xe studies“Xenology”. The Xe components include a radiogenic component ( 129Xe fromdecay of 129I with t 1/2 =17 Ma), a spontaneous fission component from 235,238 U and244Pu, a cosmic-ray spallation component, and so-called “trapped” components.Reynolds and Turner (1964) observed that the abundances of the heavy isotopes131-136 Xe in the 700-1000°C fractions released from the Renazzo CR-chondritewere akin to those produced from spontaneous fission of heavy actinides. Roweand Kuroda (1965) reported similar observations. This component was dubbedCCF-Xe (carbonaceous chondrite fission), but the ensuing debate in the late1960s about the possible presence of super-heavy elements in meteorites implied

by CCF-Xe did not last (see Anders 1988). One objection was that enrichmentsof heavy Xe isotopes were accompanied by enrichments in the lightest ones,124Xe and 126Xe, which is inconsistent with a fission origin (Manuel et al. 1972).

In their search for the carrier of the unusual Xe component in the Allendemeteorite, Lewis et al. (1975) found that, after HF-HCl treatment, essentially alltrapped Xe of the meteorite was in the residue, which comprised 0.5 mass% of

the bulk meteorite. Processing with HNO 3 removed only 6% of the residue’smass but also all trapped Xe, and a Xe component (Xe-HL), enriched in both andheavy ( 134Xe and 136Xe) and l ight ( 124Xe and 126Xe) isotopes, emerged. In Xe-HL,the isotopic anomalies in light and heavy isotopes are always correlated.

Another component with high 130Xe/ 132Xe was accidentally found whenSrinivasan and Anders (1978) tried to characterize Xe-HL in the Murchisonmeteorite (see Anders 1988). The new component was enriched in the even-numbered isotopes ( 128Xe, 130Xe, and 132Xe), and a comparison to nucleosynthesistheory indicated a signature of the slow neutron capture process ( s-process), thusit was named Xe-S. Typically, Xe-S is accompanied by a Kr component also richin s-process isotopes (Kr-S).

3. Isolation and discovery of presolar grains

Abundances of the exotic noble gas components compared to the total noblegas content are relatively low, so the abundances of the carrier minerals (what we

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now call presolar grains) of the anomalous noble gases were also expected to bequite low. For example, Lewis et al. (1975) demonstrated that the carrier of Xe-HL comprises less than 0.5% by mass of the whole Allende meteorite .

At the time the anomalous noble gas components were discovered, the carrier minerals were not known, but further studies suggested that they arecarbonaceous (Table 1; e.g., Eberhardt et al. 1981, Jungck 1982, see summary byAnders 1987, 1988). Lewis et al. (1987) succeeded in isolating diamond, thecarrier of Xe-HL. This was followed by the isolation of SiC, the carrier of Kr-S,Xe-S, and Ne-E(H) (Bernatowicz et al. 1987, Tang and Anders 1988), andgraphite, the carrier of Ne-E(L) (Amari et al. 1990).

The procedure to separate presolar grains developed by Amari et al. (1994) is

shown in Fig. 2. First, silicates, which comprise a major part of stony meteorites(~96%), are removed with HF. Part of the organic matter - aromatic polymerscollectively called reactive kerogen - is destroyed by oxidants (H 2O2, Cr 2O7

2-). Atthis point, nearly 99% of the meteorite is dissolved. From the remaining residue,

Fig. 2. The steps of the chemical isolation procedure and the resulting presolar grain size-and density-fractions from the Murchison meteorite by Amari et al. (1994).

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diamond is collected by so-called colloidalseparation. In a basic solution diamond

particles are negatively charged, repulseeach other and stay in solution, while inacidic solution they become neutral andcoagulate (Fig. 3), and settle down bycentrifugation.

Presolar graphite has similar chemical properties as other carbonaceous matter that remains in the meteoritic residue, anda density separation is applied to extract

presolar graphite from the remainingcarbonaceous matter and other mineralswith higher densities. The separate withdensity >2.2g/cm 3 are purified with HClO 4,which leaves minerals such as SiC andaluminous oxides.

Fig.3 . Nano-diamonds precipitate as acloudy white gel from acidic solution butthey completely “dissolve” in basic solution.

Photo courtesy of Roy S. Lewis.

Table 1 . Early results on presolar grain identifications ca. 1988Working designation C-α C-β a C-ε a C-δ C-θ Identified phase amorphous

carbonSiC SiC diamond amorphous

carbon Noble gas identifier Ne-E(L) b Xe-S Ne-E(H) Xe-H(L) no noble

gasesGrain size µm 0.8 – 20 c 0.03 – 20 c 0.001-

0.00250.2 - 2

Average density

(g/cm3

)

1.6–2.2 3.22 3.22 2.22 to 2.33 <2.2

Abundance d (ppm) < 2 – < 5 3 – 7 400 105δ 13C (per-mil) e +340 +1000 +1500 ~ -38 -50δ15 N (per-mil) e > +252 ~ -500 ~ -375Possible stellar source f

novae, SN, AGB? AGB AGB SN SN?

Sources : Lewis et al. 1987, Anders 1988, Tang and Anders 1988, Tang et al. 1988, Anders andZinner 1993a Originally it was thought that C- β and C- ε are two different carriers. It was believed that C- βcarries Xe-S, and that C- ε , associated with spinel because of similar density, carries Ne-E(H).,e.g., Anders (1988).

b Ne-E(L) = “pure” 22 Ne is mainly from 22 Na decay but may have some 22 Ne from AGB starsc Large grains are extremely rare.d Abundance (ppm by mass) in CM2 chondritese Values are for carrier-enriched fractions, not for the pure carrier phasef Stellar sources. AGB = asymptotic giant branch star (section 6), SN = Supernova (section 7)

Most of the grains were isolated from meteorites by progressively harsher acid dissolution of their meteorite hosts. The advantage of this method is that

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individual grains can be studied for their overall morphology and crystalstructure. The disadvantage is that other potential presolar minerals aredestroyed. Another desired piece of information is how presolar grains interactwith the meteorite host during metamorphism, because decreases in presolar grain abundances with metamorphic type are observed. Locating the grains in

situ is necessary to study grain destruction in meteorite parent bodies and toidentify their reaction products. The first in-situ detection of presolar grains inmeteorites was reported by Alexander et al. (1990).

4. Characterization techniques

The techniques used to analyze presolar grains are listed in Table 2. Grainmorphology and compositions are examined by scanning electron microscopy,

SEM. Structural information is obtained by transmission electron microscopy,TEM, and Raman spectroscopy. Thin-sections of graphite and SiC grains for TEM studies are prepared by two methods. One is to embed grains in resin and toslice them into sections with a nominal 70 nm thickness using an ultramicrotomeequipped with a diamond knife (e.g., Bernatowicz et al. 1991). Another methodis “focused ion beam lift-out”, where a ~3µm-thick Pt-layer is deposited onto agrain to protect it from subsequent sputtering by a high-energy Ga beam that“cuts” parallel trenches to make a ~100 nm thin section (e.g., Stroud et al. 2002).

Most light element isotopic data have been obtained by ion probe. Presolar oxides are a minor oxide population in meteorites, and the majority of meteoriticoxides is isotopically normal. Ion imaging techniques, pioneered by P. Hoppeand further developed by Nittler et al. (1997) are used to efficiently locate

presolar oxide grains. The ratio of the signal strengths of 16

O and18

O is used to produce isotopic “maps” of grains on a sample mount ( ∼100 ×100 µm/image).Any grains with isotope ratios significantly different than the standard then can

be located from the isotope map of the mount for further analyses.Abundance and isotopic measurements of reasonable precision for several

elements by conventional ion probe require grains of >1µm and relatively highelemental abundances. Analyses of sub-micron grains with much higher

precision are possible with the NanoSIMS. Its Cs + primary beam size can be assmall as 30 nm (as opposed to a few µm for the CAMECA IMS-3f) and achieveshigh spatial resolution. The NanoSIMS also has higher sensitivity at high-massresolution (e.g., 40 × higher than the IMS-3f for Si isotopic analysis) and a multi-detection system with four mobile and one fixed electron multipliers.

Very high precision data for the isotopic composition of heavy elements (Sr and Ba) in aggregate grain samples have been obtained by conventional thermalionization mass spectroscopy, TIMS. The sample (e.g., SiC grains) is directlyloaded onto a filament and gradually heated, which ionizes different elements atdifferent temperatures, and thus separates elements within similar mass ranges.

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Table 2. Analytical methods applied in presolar grain studiesAnalytical quantity a Method b References c

Elemental abundancesSiC (single/bulk): Mg, Al, Ca, Ti, V,

Fe, Sr, Y, Zr, Ba, CeSIMS Amari et al. 1995c

graphite (single): H, N, O Al, Si SIMS Hoppe et al. 1995diamond (bulk): Sc, Cr, Fe, Co, Ni,

Ru, Os, Ir INAA Lewis et al. 1991b

Isotopes of light elementsSiC (single/bulk): C, N, O, Al-Mg, Si,

Ca, Ti, FeSIMS Hoppe et al. 1994, Huss et al. 1997,

Amari et al. 2000agraphite (single/bulk): C, N, O, Al-Mg,

Si, Ca, TiSIMS Amari et al. 1993, Hoppe et al. 1995,

Travaglio et al. 1999diamond (bulk): C, N, MS Russell et al. 1996Si3 N4 (single): C, N, Si SIMS Nittler et al. 1995oxides (single): O, Al-Mg (Ca, Ti) SIMS Huss et al. 1994, Hutcheon et al. 1994,

Nittler et al. 1997, Choi et al. 1998,1999

silicates (single): O, Al-Mg SIMS Messenger et al. 2003, Nguyen andZinner 2004, Nagashima et al. 2004

Isotopes of heavy elementsSiC (bulk): Sr, Ba, Nd, Sm TIMS,

SIMSOtt and Begemann 1990, Zinner et al.1991, Prombo et al. 1993

Sr, Zr, Mo, Ba, Ru RIMS Nicolussi et al. 1997, 1998a,c, Savinaet al. 2003b, 2004

graphite (single): Zr, Mo RIMS Nicolussi et al. 1998bdiamond (bulk): Te TIMS Richter et al. 1998

Noble gasesSiC (bulk): He, Ne, Ar, Kr, Xe NGMS Lewis et al. 1990,1994graphite (bulk): Ne, Ar, Kr, Xe NGMS Amari et al. 1995a

SiC & graphite (single): He, Ne NGMS Nichols et al. 2005Crystal structureSiC TEM,

RamanVirag et al. 1992, Bernatowicz et al.1992, Daulton et al. 2002, 2003

graphite TEM,Raman

Zinner et al. 1995, Bernatowicz et al.1991,1996, Croat et al. 2003

diamond TEM,EELS

Bernatowicz et al. 1990, Daulton et al.1996

Multiple properties (grain size,morphology, elemental compositions)

SEM Hoppe et al. 1994, 1995

Notes. a single: individual grain analysis. – bulk: analysis of aggregate grain samples b EELS: Electron energy loss spectrometry. – INAA: Instrumental neutron activation analysis. -MS: Mass spectrometry. –NGMS: Noble gas mass spectrometry. - SEM: Scanning electronmicroscopy. - SIMS: Secondary ion mass spectrometry. - TEM: Transmission electronmicroscopy. - Raman: Raman spectroscopy. – RIMS: Resonant ionization mass spectrometryc Only a few references are listed here to provide a starting point to analytical details

Heavy element isotopes in single grains were analyzed by resonant ionizationmass spectrometry (RIMS) with the CHARISMA instrument at Argonne

National Laboratory in Chicago (e.g., Nicolussi et al. 1998a). First, a plume of

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neutral atoms and molecules is generated by laser-induced thermal desorption.Then the atoms are resonantly ionized by tuned laser beams, and subsequentlyanalyzed by a time-of-flight mass spectrometer. Although ionization schemesmust be developed for each element, an advantage is that isotopes of interest can

be analyzed without isobaric interference from other species.

5. The presolar grain “zoo”

The abundant presolar minerals in primitive meteorites are diamond, SiC,graphite, corundum, spinel, and silicates; less frequently found are Si 3 N4,hibonite, and TiO 2 (Table 3). Presolar graphite grains also often enclose smalltrace element carbide and Fe-Ni metal particles.

Table 3 . Currently known presolar mineralsMineral Characteristic size

Possible stellar source a

Discovery papers

Diamond 2 nm AGB?, SN? Lewis et al. 1987SiC 0.1 – 20 µm AGB, SN,

novaeBernatowicz et al. 1987, Tang andAnders 1988

graphite 1 – 20 µm AGB, SN Amari et al. 1990carbides in graphitemetal grains ingraphite

10 – 200 nm10 – 20 nm

AGB, SNSN

Bernatowicz et al. 1991, 1996,Croat et al. 2003, 2004

Si3 N4 0.3 – 1 µm AGB?, SN Nittler et al. 1995corundum (Al 2O3) 0.2 – 3 µm RGB, AGB,

SN?Hutcheon et al. 1994, Nittler et al.1994

spinel (MgAl 2O4) 0.2 – 3 µm RGB, AGB,SN?

Nittler et al. 1997, Choi et al. 1998

hibonite (CaAl 12O19) 0.2 – 3 µm RGB, AGB,SN?

Choi et al. 1999

TiO 2 Nittler and Alexander 1999silicates(olivine, pyroxene)

0.1 – 0.3 µm RGB, AGB,SN?

Messenger et al. 2003 (in IDPs), Nguyen and Zinner 2004 (inchondrites)

Notes. a AGB: asymptotic giant branch stars. RGB: red giant branch stars. SN: supernovae.A “?” indicates not known/uncertain.

By now, several thousand individual SiC and a few hundred graphite grainshave been analyzed. Although diamond was discovered first and is the mostabundant presolar mineral, its nm-size prevented individual grain analyses, andanalyses are restricted to samples consisting of collections of grains (Lewis et al.

1987, 1991a, Richter et al. 1998). Most of the several hundred presolar oxidegrains analyzed are corundum and spinel, and only some dozen are hibonite(Huss et al. 1994, Hutcheon et al. 1994, Nittler et al. 1994, 1997, 1998, Choi etal. 1998, 1999, Nittler and Alexander 1999, Krestina et al. 2002, Nguyen et al.2003, Zinner et al. 2003). The first presolar silicates (olivine and pyroxene) in

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meteorites were detected by Nguyen and Zinner (2004). Nagashima et al. (2004)located presolar silicates in thin sections of two primitive meteorites.

Carbonaceous presolar grains are ubiquitous in essentially all types of primitive chondritic meteorites (Table 4). The observed abundance of presolar minerals depends not only on the resistance to chemical acid processing, whichleaves abundances of diamond, graphite, and SiC essentially unaffected,depending on the specifics of the acid treatment (Amari et al. 1994), but also ongrain survival during metamorphism on meteorite parent bodies.

Abundance estimates for presolar corundum, spinel, silicates, and Si 3 N4 areonly available for a few meteorite groups. Corundum in unequilibrated Hchondrites is ~0.03 ppm and 0.13 ppm in CM chondrites, which also contain 1-2

ppm spinel and 0.002 ppm Si 3 N4. A first lower abundance limit of ~25 ppm for presolar silicates in a very primitive carbonaceous chondrite was reported by Nguyen and Zinner (2004). In addition, presolar silicates were found ininterplanetary dust particles (IDPs), where they appear to be more abundant(~5500 ppm in cluster IDPs) than in meteorites (Messenger et al. 2003).

Table 4. Abundance estimates of presolar minerals in different meteorites (ppm by mass)Chondrite group Diamond SiC GraphiteCI 940 –1400 14 10CM 400 – 740 4–14 5-6CR 400 0.6CO 300 –520 1-3 <0.15?CV-reduced 545-620 0.17-0.39 below detection limitCV-oxidized 240 –500 0.006-0.2 <0.20CH 87 0.41 0.13H 3.4 ~36 0.063L 3.4/3.7 54-64 0.008–0.08LL3.0/3.1 100 – 130 0.39–1.52EH3-4 50 –67 1.3–1.6Sources: Huss and Lewis 1995, Huss et al. 2003, Ott 1993, Zinner et al. 2003

Primitive meteorites of low petrographic type that never experienced highmetamorphic temperatures or aqueous alteration (e.g., unequilibrated ordinarychondrites) provide the best time capsules for preserving presolar grains. Earlystudies showed that the C-isotopic compositions in the most primitive ordinarychondrites (petrologic type 3) varied more than those of higher petrologic types(4-6) (Swart et al. 1983). This could suggest selective destruction of presolar grain populations in ordinary chondrites of higher-metamorphic type, which isalso indicated by noble gas measurements of diamond separates from different

types of chondrites and provides a useful measure of meteorite parent bodymetamorphism (e.g., Huss and Lewis 1994a, 1995).

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5.1. Presolar silicon carbide

Silicon carbide is the most extensively studied presolar mineral, both in the formof aggregated “bulk” samples and of single grains, for several reasons. First, theseparation procedure of SiC is less complicated than that of graphite. Second,SiC is present in several classes of meteorites (Table 4). Third, although mostSiC grains in meteorites are sub-micron in size, some grains larger than 1 µm aresuitable for correlated isotopic analyses in single grains. Fourth, trace elementconcentrations are high enough to make elemental abundance and isotopicmeasurements with reasonable precision.

The average size of SiC grains is ~0.5 µ m in the Murchison meteorite, withsizes ranging up to 20 µ m, but grains >10 µ m are extremely rare. The sizedistribution of SiC in different types of meteorites varies, e.g., the Murchisonchondrite typically has larger-size SiC on average than other meteorites, for unknown reasons (Amari et al. 1994, Huss et al. 1997, Russell et al. 1997).

The morphology of most SiC grains extracted from meteorites by chemical procedures shows euhedral shapes with more or less pitted surfaces (Fig. 4a,b).

In order to examine “pristine” SiC grains not subjected to chemicals,Bernatowicz et al. (2003) dispersed matrix material excavated from the interior of the Murchison meteorite onto polished graphite planchets and examined them

by SEM. The pristine grains have less surface pits than grains isolated bychemical procedures, indicating that etching of surface defect structures occursduring chemical isolation. Of the 81 pristine grains studied by Bernatowicz et al.(2003), ~60% are coated with an amorphous, possibly organic phase. No

Fig. 4. Secondary electron images of SiC grains from the Murchison meteorite. Larger grainssuch as these shown are relatively rare. Scale bars are 1 µm. (a) The pitted surface structure iscommon for SiC grains, and most likely due to the harsh chemical treatments during theextraction from meteorites. The 12C/13C ratio of this grain is 55 (cf. solar = 89). (b) A SiC grainwith a smooth surface. The 12C/13C ratio of this grain is 39.

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differences in morphology other than those caused by sample extraction procedures have been observed among SiC grains.

Synthetic SiC has several hundred different crystallographic modifications but presolar SiC apparently only occurs in the cubic 3C and hexagonal 2H modification andintergrowths of these two (Daulton et al. 2002, 2003). This limited polytype distributionin presolar SiC suggests condensation of SiC at relatively low total pressures incircumstellar shells (Daulton et al. 2002, 2003).

5.1.1. Chemical and isotopic compositions of “bulk” SiC aggregates

The analyses of bulk samples have the advantage that data with high precisioncan be obtained. This provides well-defined average properties for whole suitesof separates which allows one to examine systematic differences among thevarious fractions. On the other hand, there is the potential problem thataggregates may contain some contaminant, which can hamper the interpretation

of bulk elemental abundances. However, any isotopic anomalies would only bediluted by contamination with normal material, and large overall isotopicanomalies can still be detected.

Silicon carbide is the carrier of Ne-E(H), Kr-S and Xe-S. Lewis et al. (1990,1994) analyzed all noble gases in size-sorted SiC fractions (average size: 0.38-4.6 µm) extracted from the Murchison meteorite. They modeled the measurednoble gas compositions by a mixture of two-components, which they designated

Fig. 5. Noble gas components in aggregates of SiC grains (Lewis et al. 1990, 1994, Ott 2002)normalized to solar isotopic composition (Wieler 2002). Isotope ratios are further normalized to20 Ne, 36Ar, 84Kr, and 132Xe, respectively. The dotted line shows solar composition.

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“N” component for “ normal”, and “G” component for “A GB” (Fig. 5). The Gcomponent of Kr was found to be similar to the theoretical composition fromnucleosynthesis by the s-process in AGB stars (Gallino et al. 1990, 1997), hencethe name. Initially the terms “G” and “N” component were only used for noblegases but are also in use for other elements. The isotopic compositions of Sr (Podosek et al. 2004) and Ba (Zinner et al. 1991, Prombo et al. 1993) wereanalyzed for the same size-sorted SiC fractions (but for different suites) by TIMSand SIMS. In all cases, the abundant G component indicated that a large fractionof presolar SiC grains came from stars in which the s-process responsible for theG component operates.

5.1.2. Individual SiC grains: Clues to SiC sub-populations

Isotopic and trace element analyses of single grains by ion probe revealed the presence of different types of SiC. The isotopic compositions of C, N, and Si leadto five major SiC subtypes called mainstream, A+B, X, Y, and Z (Table 5, Figs.6 and 7). The number of grains plotted in Figs. 6 and 7 does not represent their true distribution among the sub-populations because certain grains were

preferentially searched for by ion imaging and then studied. The “true”distribution (by number) is noted in the legends of Figs. 6 and 7, and in Table 5.

12 C/ 13 C1 10 100 1000 10000

1 4 N / 1

5 N

1

10

100

1000

10000

U

SiCmainstream ~93%

A+B 4-5%

X ~ 1%

Y ~1%

Z ~1%

nova

Si3N4

Fig. 6. SiC grains fall into different populations based on their C- and N-isotope ratios(Alexander 1993, Amari et al. 2001a-c, Hoppe et al. 1994, 1997,2000, Huss et al. 1997, Linet al. 2002, Nittler et al. 1995). For comparison, stellar data are plotted with error bars andtheir N-isotopic ratios are typically lower limits (Wannier et al. 1991, Querci and Querci1970, Olsen and Richter 1979). The dotted lines indicate solar isotope ratios.

15



About 93% (by number) of presolar SiC are mainstream grains , with lower 12C/13C and higher 14 N/ 15 N than the respective terrestrial reference ratios of 12C/13C=89 and 14 N/ 15 N=272. Their Si-isotopic composition is slightly 29Si- and30Si-rich ( 29Si/28Si and 30Si/ 28Si are up to 1.2×solar). In the three Si-isotope plot,mainstream grains define a line with δ29Si = –15.9 + 1.31 δ30Si (Lugaro et al.1999). The fit parameters vary depending on the number of points included in theregressions, e.g., Hoppe et al. (1994) find δ29Si = –15.7 + 1.34 δ30Si, and Nittler and Alexander (2003) obtained δ29Si = –18.3(±0.6) +1.35(±0.01) δ30Si .

Grains with 12C/13C < 10 and 14 N/ 15 N = 40 – 12,000 are called A+B grains (Amari et al. 2001a). Originally, it was thought that two distinct populations “A”

and “B” existed, but these two belong to the same continuum spanned by C- and N-isotopes (Fig. 6). The Si-isotopes of A+B grains are similar to those of mainstream grains (Fig. 7). In the three Si-isotope plot A+B grains define a linewith δ29Si = –34.1(±1.6) + 1.68(±0.03) δ30Si (Amari et al. 2000b) with a small

off-set in slope compared to the mainstream grains. The A+B grains are thesecond largest presolar SiC population and constitute 3-4% of all SiC grains.

δ 30 Si-800 -600 -400 -200 0 200 400 600 800

δ 2 9 S i

-800

-600

-400

-200

0

200

400

600

SiCmainstream ~93%A+B 4-5%X ~ 1%Y ~1%Z ~1%

novaSi3N4

Fig. 7. The Si-isotopes for presolar SiC grains (references as in Fig. 6) and stars (symbolswith error bars) are given in the δ−notation which describes the deviation of an isotope ratio(i N / j N ) of a sample from the (terrestrial) standard ratio in per-mil: δ i N (‰) = [( i N / j N )sample

/( i N / j N )standard – 1] × 1000. The grains fall into distinct populations. The triangles show twodeterminations for the C-star IRC+10°216 (Cernicharo et al. 1986, Kahane et al. 1988). Theother stellar data (circles) are for O-rich M-giants (Tsuji et al. 1994).

The SiC grains of type X , ~1% of all SiC, have higher 12C/13C and lower 14 N/ 15 N than the respective solar ratios. The X grains have low δ29Si and δ30Si

16



values, and 28Si excesses reach up to 5×solar (Amari et al. 1992, Hoppe et al.2000, Amari and Zinner 1997). Another one percent of all SiC grains have12C/13C >100 and 14 N/ 15 N above the solar ratio (Amari et al. 2001b, Hoppe et al.1994). These Y grains appear to be 12C-rich mainstream grains but their 30Si/ 28Siratios are slightly larger than in mainstream grains, which merits placing theminto a separate group. Up to 3% of all SiC grains, particularly among smaller sizegrain fractions, are Z grains . Their 12C/13C and 14 N/ 15 N ratios are similar to thoseof mainstream grains, but Z grains have large 30Si excesses (Alexander 1993,Hoppe et al. 1997). Only a few nova SiC grains, with 12C/13C = 4-9, and 14 N/ 15 N= 5-20, are known (Amari et al. 2001c, José et al. 2004, Nittler and Hoppe2004a,b). Most nova grains have close-to-solar 29Si/ 28Si but 30Si excesses.

Table 5. Some characteristics of presolar silicon carbide populationsDesignation Mainstream X Y Z A+B a Nova

Crystal type 3C, 2Hb

3C, 2Hb

3C, 2Hb

3C, 2Hb

3C, 2Hb

3C, 2H?

b Heavy traceelements c

~10-20× c highlydepleted

~ 10× c NA solar or 10-20× c

NA

12C/13C 10 – 100 20 – 7000 140 – 260 8 – 180 < 3.5 (A)3.5 – 10(B)

< 10

14 N/ 15 N 50 – 2×104 10 – 180 400 – 5000

1100 – 1.9×10 4

40–1.2×10 4 < 20

29Si/ 28Si c 0.95-1.20× 28Si-rich 0.95-1.15×

≈solar 1.20× ≈solar

30Si/ 28Si c 0.95-1.14× 28Si-rich 30Si-rich 30Si-rich 1.13× 30Si-rich26Al/ 27Al 10 -3 to 10 -4 0.02 to 0.6 similar to

MSsimilar to

MS<0.06 up to 0.4

Other isotopicmarkers c

excess in46Ti, 49Ti,

50Tiover 48Ti

44Caexcess

41K excess


50Tiover 48Ti


49Tiover 48Ti


50Tiover 48Ti

47Ti-rich

22 Ne d yes NA NA NA NA NAAbundance 87-94% 1% 1 – 2% 0 – 3% 2 – 5% << 1%Sources : Amari et al. 2001a,b, Hoppe and Ott 1997, Hoppe and Zinner 2000, Nittler and Hoppe2004a,b,Ott 2003, Zinner 1998a Group A and B grains were initially separated but later found to form a continuum in

composition. b cubic 3C, hexagonal 2H; Daulton et al. (2002, 2003).c Abundance compared to solar composition.d 22 Ne = Ne-E(H) = Ne(G); and NA: not analyzed.

Many SiC grains have 26Mg/ 24Mg larger than the solar ratio but solar 25Mg/ 24Mg within 10% (Amari et al. 1992, 2001a-c, Hoppe et al. 1994, 2000,Huss et al. 1997). Magnesium in some X grains is almost pure 26Mg, and 26Mgexcesses are most likely from in-situ decay of 26Al (t 1/2 = 7.3×10 5 a) that wasincorporated into grains at their stellar sources. The X grains have 26Al/ 27Al of upto ~0.6 (Fig. 8), whereas ratios in A+B and mainstream grains typically do not

17



12C/ 13C

1 10 100 1000 10000

2 6 A l / 2

7 A l

10 -5

10 -4

10 -3

10 -2

10 -1

10 0

graphite

Si3N4

SiC - A+B

SiC-mainstr.

SiC - X

SiC - Y

SiC - nova

exceed 0.01 (Hoppe et al. 1994, Amari et al. 2001a). Isotopic compositions werealso measured for Ca and Ti (Ireland et al. 1991, Amari et al. 1992, 2001a,b,Hoppe et al. 1994, 1996, 2000, Nittler et al. 1996, Alexander and Nittler 1999,Hoppe and Besmehn 2002, Besmehn and Hoppe 2003) and Fe, Zr, Mo, Sr, Ba,and Ru (e.g., Nicolussi et al. 1997, 1998a, 1998c, Pellin et al. 2000a,b, Davis etal. 2002, Savina et al. 2003a,b, 2004).

Fig. 8. Inferred 26Al/27Al ratios vs. 12C/13C ratios in SiC and low-density graphite grains.The SiC type X and graphite grains have the largest 26Al/ 27Al. See text for data sources.

5.1.3. Trace elements in individual SiC grains

Silicon carbide grains contain several trace elements, some of them inconsiderable amounts. Magnesium concentrations are typically around 100 ppmand Al abundances can reach several mass-percent (e.g., Hoppe et al. 1994).

Nitrogen, probably substituting for carbon in the SiC lattice, shows relativelyhigh concentrations so the N- isotopic ratios can be analyzed with reasonable

precision. However, determination of the absolute concentration of N is difficultas carbon must be present to produce CN - which is used to analyze N by the ion

probe (Zinner et al. 1989). In addition to Al and Mg, concentrations of Ca, Ti, V,Fe, Sr, Y, Zr, Nb, Ba, Ce, and Nd were measured in 60 SiC grains (average size:4.6 µ m) and in three size-sorted SiC aggregates of 0.49-0.81 µ m (Amari et al.

1995c). The general problem in determining multiple trace element abundancesin presolar grains is that grains are partially consumed during the measurements.The first measurements are generally for C, N, and Si isotopic ratios to identify

18



the grain’s membership among the SiC populations. After this, only initiallylarge grains have enough mass left for additional trace element analyses.

Trace element measurements of 34 mainstream grains define at least 8-9different normalized abundance patterns (Fig. 9), and only 4 of these grains hadunique abundances (Amari et al. 1995c). In mainstream grains, element/Si ratiosof elements heavier than iron (Y, Zr, Nb, Ba, La, Ce, and Nd) show up to 35

Fig. 9. Trace element abundances in individual SiC grains normalized to solar abundances andSi. Relative depletions or enrichments in Sr and Ba are most notable. Mainstream grains areshown by black symbols, open symbols are for A+B grains, and type Y grains are shown ingrey. Abundances in two X type SiC grains are shown in a separate panel. The letters marked

by a star refer to the original groupings in Fig. 1 by Amari et al. (1995c).

19



times the solar ratio. These enrichments most likely reflect the composition at thestellar sources. However, the source composition can be fractionated according tovolatility during trace element condensation into SiC (Lodders and Fegley 1995).In Fig. 9, the s-process elements Sr and Ba should be as abundant as their neighbors, but they are depleted in many grains because they are more volatile.

The enrichment of s-process elements in Y grains (Fig. 9) supports their closerelationship to mainstream SiC. Trace elements have not yet been reported for Ztype grains. Most of the trace element abundances measured in the two X grainsare very low and data for Zr and Sr are upper limits. The trace elements of twentyA+B SiC grains (Fig. 9) give 4 different abundance patterns (Amari et al. 1995c,2001a). This led to the surprising conclusion that A+B grains require at least twotypes of stellar sources with different characteristic heavy element abundances,which was not indicated by their isotopic compositions. In one of the patterns, allelements plot near the solar abundance ratio, and only Sr shows a relativedepletion. The other 3 patterns show higher relative abundances of the heavyelements and various volatility-related fractionations, which makes these A+Bgrains more similar to mainstream grains.

5.2. Presolar silicon nitride

The few known Si 3 N4 grains share many properties of the SiC X grains (Figs. 6-8). Information about these rare grains is given by Amari et al. (1992), Besmehnand Hoppe (2003), Lin et al. (2002), Nittler and Alexander (1998), and Nittler etal. (1995).

5.3. Presolar graphite

Presolar graphite is present only in the least thermally altered primitivemeteorites (Huss and Lewis 1995, Table 4). It is not as chemically resistant asSiC or the oxides corundum and spinel, and its isolation is complicated becauseother carbonaceous compounds with similar chemical and physical properties are

present in primitive meteorites. The separation of graphite, which is achieved bya combination of mild oxidation and density separation, is far more elaboratethan that of other presolar grains, and essentially all studies of presolar graphitehave been performed on four graphite-rich fractions extracted from theMurchison meteorite (Amari et al. 1994). Some properties of these four fractionsare given in Table 6.

Some structural, elemental, and isotopic features of presolar graphite grains

vary with density. On average, the low-density separates contain more largegrains (Hoppe et al. 1995). Amari et al. (1995a) and Hoppe et al. (1995) describecorrelations between density and isotopic compositions of the noble gases andcarbon. This is in marked contrast to SiC grains, where observed isotopic featuresin Kr, N, Sr, and Ba depend only on grain size. Graphite grains are much larger,

20



typically >1 µm, but ranging up to 20 µm, than most sub-micron SiC grains.Investigation by SEM and TEM show two morphological types (Fig. 10) -dubbed “cauliflower” and “onion”- (Bernatowicz et al. 1991, 1996, Hoppe et al.1995, Bernatowicz and Cowsik 1997). Grains of the onion type have aconcentric-layered structure (reminiscent to that of hailstone, although on adifferent absolute scale) of well-graphitized carbon and a core of randomlyoriented, fine-grained crystalline aggregates (Bernatowicz et al. 1996).

Table 6. Some characteristics of presolar graphite density fractions a Designation KE1, KE3 KFA1 KFB1 KFC1Density (g/cm 3) 1.6 – 2.05 2.05 – 2.10 2.10 –2.15 2.15 – 2.20Morphology mainly cauliflower mainly onion12C/13C 3.6-7223 3.0-2146 3.8-3377 2.1-406414 N/ 15 N mostly solar

(28-306)mostly solar (123-398)

mostly solar (153-315)

mostly solar b

18O/16O up to 184 × solar up to 6.6 × solar ∼solar ∼solar 29Si/ 28Si 0.63-2.3 × solar 0.54-1.57 ×

solar ± solar c 0.8-1.3 × solar

c

30Si/ 28Si 0.46-1.94 × solar 0.39-1.4 × solar ± solar c ± solar c,d 26Al/ 27Al up to 0.146 up to 0.138 up to 0.086Major sources: Hoppe et al. (1995), Travaglio et al. (1999)a Amari et al. (1994).

b One grain in KFC1 with 14 N/ 15 N = 730±64.c Very large uncertaintiesd One nova grain with 30Si/ 28Si = 1.76 × solar.

The cauliflower type grains appear to be aggregates of small grains andconsist entirely of turbostratic graphite (i.e., graphite with contorted layers nothaving long-range continuity, Bernatowicz et al. 1991, 1996). The onion typegraphite is more abundant in high-density fractions, and the cauliflower type ismore abundant in low-density fractions.

Fig. 10. Presolar graphite grains show two morphologies. (a) A graphite grain of the“onion” type, with a layered surface structure. (b) A graphite grain of the “cauliflower”type, which appears as aggregates of small grains.

21



Fig. 11 shows the C- and N-isotopic composition of individual graphite grainsfrom the four density fractions. Overall, the 12C/13C ratios range from 2.1 to7223, demonstrating their presolar origin. Grains with 12C/13C higher than thesolar ratio are most abundant (~75%) in the highest density separate KFC1. Incontrast, the isotopic composition of

N, one of the most abundant traceelements in graphite, is bewilderinglysolar in many grains. This has beeninterpreted that at least part of theindigenous N in graphite equilibratedwith air (Hoppe et al 1995) Thegraphite from the lowest densityseparate (KE3) is a possibleexception because there is a slightdecrease in 14 N/ 15 N with increasing12C/13C (Fig. 11). An outstandingfeature of graphite grains is the

presence of almost pure 22 Ne (=Ne-E(L)). If N isotopically equilibratedwith air the noble gases should alsoequilibrate, thus erasing any isotopicanomalies. In particular, Ne should

be affected because it has adsorption properties similar to N, but Ne isanomalous. This compounds themystery of the ‘normal’ N-isotopiccomposition and requires moreinvestigation. Abundances of H, O,

N, Al, and Si in individual graphitesare typically about one-mass percent(Hoppe et al. 1995). With increasingdensity of the grain fractions, theconcentrations of H, N, O, and Sidecrease. Isotopic compositions of O,Mg, Al, Si, Ca, Ti, Zr and Mo in individual graphite grains from the KE3, KFA1and KFC1 density fractions were measured by e.g., Hoppe et al. (1995), Amari etal. (1995b, 1996), Nicolussi et al. (1998b), and Travaglio et al. (1999).

Fig. 11. Carbon and N-isotopes in individualgraphite grains from 4 density fractions.Density increases in alphabetical order (E<FA<FB<FC).

The Si-isotopes in graphite grains (Fig. 12) are similar to those observed for the different SiC populations. Many grains from the KE3 separate are rich in 28Si,like the SiC X grains (Fig. 7). This kinship also follows from the 26Al/ 27Al ratios(Fig. 8). A significant fraction of grains from the low-density fractions KE3(Travaglio et al. 1999) and KFA1 (Amari, unpublished) appear normal in 17O/16O

22



but show large enrichments in 18O/16O relative to solar (Fig. 14). Thiscomposition is remarkably different from that of presolar oxides.

Fig.12. Silicon isotopes in presolar graphite grains from the low-density fractions definearrays similar to those given by presolar SiC (Fig. 7). Note that the measurements may haverelatively large uncertainties, see e.g., Hoppe et al. (1995).

The KE3 grains, and a few of the KFA1 grains are believed to have formed insupernovae, whereas the high-density KFC1 grains most likely formed in low-metallicity (i.e., less metal-rich than the sun) AGB stars. In addition, novae canaccount for isotopic features of a few grains.

There are no equivalent designations for the different graphite populations asfor the mainstream, A+B, X, Y, and Z populations of SiC. It is still unknownhow many presolar graphite populations exist and how many types of stars arerequired to account for the presolar graphite. The C-isotope ratios alone are notdiagnostic enough to distinguish the possible stellar sources, and the isotopiccompositions of other elements are needed to identify the graphite grains’ parentstars. However, with the exception of the KE3 grains, the low trace elementcontents make isotopic analyses challenging.

5.4. Sub-grains in graphite and SiC

Nanometer size sub-grains hidden in presolar graphite and SiC grains made their debut during TEM studies by Bernatowicz et al. (1991, 1992). A SiC grain about5.3 µ m in size of clear presolar origin ( 12C/13C = 51.6±0.4) contained several sub-grains ranging in size from 10-70 nm. Diffraction patterns and energy-dispersive

23



X-ray spectrometry confirmed that theyare TiC. An epitaxial relationship

between the TiC and the SiC hostsuggests that SiC and TiC either grewsimultaneously, or that TiC exsolvedfrom SiC (Bernatowicz et al. 1992).More extensive studies have been

performed on inclusions in graphitegrains from the lowest density separateKE3 (Croat et al. 2003) and the highestdensity separate KFC1 (Bernatowicz etal. 1996). The tiny sub-grains in graphiteare shown in Fig. 13. The sub-grains inKFC1 graphites are smaller (5-200 nm)than those in KE3 graphites, whichcontain ~10 to ~400 nm-size TiC grains.

The variable V/Ti ratios of 0.07–0.2in sub-grains indicate an independentorigin from the host graphites. In contrast

to sub-grains in SiC, those in graphite show no crystallographic (epitaxial)relationship with the host graphite, suggesting that they formed prior to graphiteand were randomly incorporated into graphite grains later. About 30% of the TiCgrains have partially amorphous rims (3 to 15 nm thick) that could be the resultof atom bombardment when the grains were adrift before they were embeddedinto the growing graphite grains (Croat et al. 2003).

Fig. 13 . A TEM image of two TiC sub-grains in a slice of a presolar KE3 graphitegrain. Photo courtesy of K. Croat and T.Bernatowicz.

Many graphite grains from the high-density KFC1 fraction contain sub-grainsconsisting of refractory carbides with compositions ranging from nearly pure TiCto nearly pure Zr-Mo carbides (Bernatowicz et al. 1996; Croat et al. 2004). Inseveral graphite grains, a carbide particle is located in the center of the grains,indicating that the carbide served as a nucleation site (see Figure 13 and theFigure 7(a) in Bernatowicz et al. 1996). However, there are also graphite grainswith sub-grains of TiC, iron-nickel metal, and cohenite, which contain no Mo-Zr carbides. This points to different stellar sources, possibly supernovae, for thesegraphite grains (Croat et al. 2003).

5.5. Presolar oxide grains: Corundum, spinel, hibonite

Oxide grains are resistant to chemicals used to isolate carbonaceous presolar

grains, and they are also concentrated in SiC-rich residues. The problem of identifying presolar oxides is that the vast majority of oxides in meteoriticresidues is isotopically normal because of its solar system origin. For example, inthe Tieschitz (H3.6) chondrite presolar Al 2O3 is estimated to be 0.03 ppm (Nittler

24



et al. 1997), and in CM2 chondrites less than 1% of all corundum grains areexpected to be presolar (Zinner etal. 2003), so the ion imagingtechnique is valuable in locating

presolar oxides among the normalones.

18 O/ 16 O10 -5 10 -4 10 -3 10 -2 10 -1

1 7 O / 1

6 O

10 -4

10 -3

10-2

.

oxide grainsgraphite

RGB starsAGB stars

PN

group III

group II g r o u p I

V

g r o u p

I

The first presolar oxides werefound during ion probe studies onindividual grains (Huss et al. 1994,Hutcheon et al. 1994). Most

presolar oxides were found by ionimaging and subsequently analyzed

by conventional ion probe in high-mass-resolution mode. The oxidesare mainly corundum and spinel(Huss et al. 1994, Hutcheon et al.1994, Nittler et al. 1994, 1997,1998, Nittler and Alexander 1999,Choi et al. 1998, 1999, Krestina etal. 2002), a few hibonite grains(Choi et al. 1999, Krestina et al.2002), and one TiO 2 grain (Nittler and Alexander 1999).

Oxide Grains

18 O/16 O10 -5 10 -4 10 -3 10 -2

1 7 O / 1

6 O

10 -4

10 -3

10 -2

group III

group II

g r o u p I V

g

r o u p

I

> 0.03

1e-4 - 0.03

<1e-4

26Al/27Al

Corundum was the dominant presolar oxide known until Zinner et al. (2003) analyzed individual

Fig. 14. The O-isotopes in presolar oxideand silicate grains define 4 groups (seetext for references). Also shown are datafor graphite grains from the low-densityfractions KE3 (Travaglio et al. 1999) andKFA1 (Amari, unpublished); and

planetary nebulae (PN), RGB, and AGBstars (Harris & Lambert 1984, Harris et al.1985, 1987, 1988, Kahane et al. 1992,Lambert et al. 1986, Smith & Lambert1990, Wannier and Sahai 1987). Dottedlines indicate solar ratios.Fig. 15. Same as Fig. 14, but symbolcolors indicate the 26Al/ 27Al content of

oxide grains. Black symbols show grainswith 26Al/27Al > 0.03, grey with 0.03 <26Al/ 27Al < 10 -4, and white with 26Al/ 27Al<10 -4 and grains for which only upper limits have been determined.

25



(Nittler, private communication) but the spatial resolution and sensitivity of theinstrument were not high enough to identify sub-micron presolar silicate grains.

The first identification of six presolar silicates of 0.3–0.9 µ m in diameter wasin IDPs by O-isotope ion imaging with the NanoSIMS (Messenger et al. 2003).One of them is forsterite and two are amorphous silicates called GEMS (glasswith embedded metal and sulfides). Using the same technique, Nguyen andZinner (2004) identified 9 presolar silicate grains ranging between 0.2 and 0.6 µ min the size-separated disaggregated matrix from the Acfer 094 chondrite, whichhas experienced minimal thermal or aqueous processing. The grains aretentatively identified (because of interference from underlying or surroundinggrains) as pyroxene, olivine, and Al-rich silicate. Also employing advancedtechniques for high spatial resolution, Nagashima et al. (2004) identified oneolivine and five mineralogically-uncharacterized presolar silicate grains in thinsections of the Acfer 094 and NWA530 meteorites. The available O-isotope dataon presolar silicates from IDPs and meteorites are similar to those of presolar oxides (Fig. 14).

5.7. Presolar diamond

Diamond was the first presolar mineral identified in meteorites (Lewis et al.1987), and it has the highest relative abundance among carbonaceous presolar grains. Still, it remains the least understood, mainly because the diamond grainsare only 1-3 Ǻngstroms in size (Lewis et al. 1987, Daulton et al. 1996), whichmakes them too small for individual analysis. The presolar diamond separatesobtained by chemical acid treatment are not 100% pure carbon and are of lower density than normal crystalline diamond (3.51 g/cm 3). The diamonds contain N

and O, probably in chemical functional groups on their surfaces as inferred frominfrared spectra (Lewis et al. 1989, Mutschke et al. 1995, Andersen et al. 1998,Jones et al. 2004). Diamonds from ordinary-, enstatite-, and CV-chondritescontain N ranging from 2200 to 42100 ppm (by mass) and those from CI- andCM-chondrites have 7500 to 9200 ppm (Russell et al.1991). In addition, diamondseparates contain noble gases, H enriched in deuterium up to 250-340‰ (Careyet al. 1987, Lewis et al. 1989), and some trace elements (Lewis et al. 1991b).

The 12C/13C ratio of 92-93 of bulk diamond in primitive meteorites (Lewis etal. 1987, 1989, Russell et al. 1991, 1996) is surprisingly close to the solar ratio of 89, considering the wide range in C-isotopic compositions of presolar SiC andgraphite grains. The δ15 N in diamonds ranges from –330‰ to –350‰ (Lewis etal. 1987, Russell et al. 1996). This range is below the values observed in various

types of meteorites (~–100 to +200‰) but is close to the upper limit of <–240‰for solar wind implanted into the lunar regolith, and to the N-isotopiccomposition of Jupiter ( δ15 N = -374‰, Owen et al. 2001), which is probably

27



more representative for the solar system than that of the terrestrial atmosphere,which is influenced by atmospheric escape processes.

The similarity of C- and N-isotopes of presolar diamonds and the solar systemsupports the idea that a large diamond fraction originated within the solar systemand is not presolar. However, the complexity of the carbon release patterns and

C/N ratios observed during stepwise heating by Russell et al. (1996) suggests the presence of more than one type of diamond. It is also possible that another phaseremained in the diamond separate after chemical and physical isolation.

Fig. 16. Noble gas isotopic composition in presolar diamonds. Data are from Huss andLewis 1994a (see also Ott 2002) normalized to the solar isotopic composition from Wieler (2002). The isotope ratios are further normalized to 20 Ne, 36Ar, 84Kr, and 132Xe,respectively. The dotted line indicates solar composition.

Detailed studies show that at least two clearly resolved noble gas components,called “P3” and “HL”, reside in meteoritic nanodiamonds (Fig. 16). In addition, acomponent called “P6”, less clearly resolved, is released at the highesttemperatures from diamond separates (Tang et al. 1988, Tang and Anders 1988,Huss and Lewis 1994b, Ott 2002). The first component, “P3” is released around500°C. Except for Ne, the isotopic ratios in P3 are close to solar, without anymajor fractionations among individual isotopes of a given element (Fig. 16). Thesecond component, HL, released around 1300°C, shows stronger fractionationsrelative to solar. Its most prominent feature is the relative enrichment in the light

and heavy Xe-isotopes, the Xe-HL. The Kr isotopic ratios increase withincreasing mass number. The 38Ar/ 36Ar is higher than solar, and the Ne isotopiccomposition is probably a mixture of “true” Ne-HL and Ne from the thirdcomponent “P6”, which is not well characterized (Huss and Lewis 1994b, Ott

28



2002). In contrast to Ne in the P3 component, the solar-normalized 21 Ne/ 20 Neratio in the high temperature fraction is larger than the 22 Ne/ 20 Ne ratio (Fig. 16).

The low temperature component P3 is mainly seen in diamond separates fromCI- and CM-chondrites and is only marginally present in CV-, ordinary-, andenstatite chondrites (Huss and Lewis 1994b, Ott 2002). This suggests preferredremoval of the P3 component during thermal metamorphism, and its presence or absence is used to estimate the degree of parent body metamorphism (Huss andLewis 1994a, 1995).

The two major noble gas components suggest that at least two different populations of diamond exist, but the reason why one diamond population is lessthermally stable than the other is not entirely clear. However, the Nconcentrations in the diamonds of CI- and CM-chondrites are also higher than indiamond from CV-, ordinary-, and enstatite chondrites. If the diamonds carryingthe low temperature P3 component have more N as impurity in solid solution,their crystal structure is plausibly less stable against thermal destruction.

The tracer of presolar diamond, Xe-HL, indicates a supernova origin of diamonds, and isotope anomalies in other heavy elements, i.e., Te and Pd(Richter et al. 1998, Maas et al. 2001) also point to supernova nucleosynthesis.On the other hand, the C- and N-isotopic signatures indicate that not all presolar diamonds originate from supernovae and that the supernova contribution to thediamonds is probably not very large. A huge isotopic anomaly in Xe-HL within asmall grain fraction can easily dominate the overall observed Xe-isotopiccomposition but not the C- and N-isotopic composition. Carbon (and likely N) is

present in all diamonds grains, and the C and N signatures from a small fractionof supernova grains could easily be masked by those from more abundant grains.In contrast, the noble gases in the presolar diamond fraction could be dominated

by a small population of gas-rich diamond.There is no shortage of suggested origins of presolar diamonds, which

remains enigmatic. These include an origin in supernovae (Clayton 1989,Clayton et al. 1995), within the solar system itself (Dai et al. 2002), novae(Clayton et al. 1995), C-rich giant stars (Clayton 1975, Lewis et al. 1987,Andersen et al. 1998), Wolf-Rayet stars (Tielens 1990, Arnould et al. 1997), and

binary star scenarios, where mass from a carbon-rich giant star is transferred ontoa white dwarf, which then explodes as a type Ia supernova (Jørgensen 1988).

6. Giant stars and their grains

The majority of known presolar grains apparently came from giant stars, whichare steady dust contributors to the ISM. We describe some of the evolution of giant stars because their nucleosynthetic products condense into grains in thestars’ circumstellar shells. One should keep in mind that the broad theoretical

picture of nucleosynthesis and evolution of giant stars is consistent with

29



astronomical observations and measurements of presolar grains, but that manydetails are not yet completely understood.

6.1. Chemical and isotopic markers of giant star evolution

Observationally the different stages of stellar evolution are tracked in theHertzsprung-Russell (HR) diagram (Fig. 17). It shows the difference in the fluxof star light between certain optical wavelengths (the “color difference”), versusabsolute magnitude. The color difference is a measure of stellar effective(photospheric) temperatures. The absolute magnitude is the energy flux at a given

wavelength normalized to astandard distance to takedistance-related dimming of starlight into account, and it

provides a measure for stellar brightness or luminosity. In theHR diagram, cool stars plot onthe right side, and the brighteststars plot at the top. The diagonal

band of stars is the “mainsequence” and consists of dwarf stars (including the Sun) withincreasing mass from bottom totop. The branches off the mainsequence are occupied by giantstars.

Red giants are in late stages of stellar evolution, and haveevolved from dwarfs with mainsequence masses of ~1 to ~8M ? (the “ ? ” refers to solar units,e.g., 1 M ? = 1 solar mass). Theevolution of low- to intermediatemass stars is sketched in Fig. 18,and more details are described by

Iben and Renzini (1983), Iben(1991), Busso et al. (1995), Arnett(1996), Gallino et al. (1997),

Wallerstein et al. (1997), Lattanzioand Forestini (1999), and Pagel (1997).

Fig. 17. The Hertzsprung-Russell diagram separatesdwarf stars on the main-sequence from stars on thegiant star branches. Data are for 11760 stars within100 parsec with distance measurement uncertainties<5% (HIPPARCHOS catalogue).

Stars spend most of their life on the main-sequence where they burn H to Hein their cores. This hydrostatic burning lasts ~10 Ga in a 1 M ? star but is shorter

30



in more massive stars. After exhaustion of H in the core, the remaining He-richcore (the “He-core”) contracts, and rising temperatures ignite H-burning in a thinshell around the He-core. The star’s radius expands to sizes of 10-50 R ? (e.g.,Jura 1999) and the star becomes more luminous as it moves onto the “red giant

branch” (RGB).Hydrogen-burning proceeds through the proton-proton (pp)-chain, and, more

importantly in higher mass stars ( á 1.1 M ? ) with higher core temperatures,through the carbon-nitrogen-(CN) or carbon-nitrogen-oxygen-(CNO) cycle,where isotopes of C,N, and O act as catalysts to convert four protons into He. Inthis process, the sum of C, N, and O nuclei remains about constant but itincreases the N abundance and decreases the C and O abundances relative to themain-sequence composition (usually taken as solar if not stated otherwise).Steady-state (or “equilibrium”) values of the CNO-cycle are 12C/13C ~3.5, and,depending on temperature, 14 N/ 15 N ~30,000 (at low T) to 14 N/ 15 N <0.1 (at highT).

Table 8. Some properties of giant stars of < 9 M ? and the SunSun M M S C J

Status main-sequence

red giant branch, RGB

asymptotic giant branch, AGB possiblyRGB?

C/O 0.5 < 0.5 ~0.5–0.7 0.6–1 1.2 1.1 s-processelements a

solar solar > solar > solar > solar ~solar

12C/13C 89 6 - 20 10-30 50-70 30-80 3 – < 1014 N/ 15 N 272 ? ? ? >500

(4–12)×10 3 ~70, ~150

17O/16O 3.78×10 -4

b (0.91–6.25)

×10 -3 (0.9–6.3)

×10 -3 (0.33–1)

×10 -3 (0.18–2.4)

×10 -3 (0.24–1.4)

×10 -3 18O/16O 2.01×10 -3

b

2.0×10 -3 (0.21–1)

×10 -3

(0.2–1)

×10 -3

(0.42–1.4)

×10 -3

6.35×10 -4

Isotopicchanges c

- +13C, + 14 N,+17O

+12C, + 22 Ne +13C, + 14 N

Major sources : Dominy et al. (1986), Harris and Lambert (1984), Harris et al. (1985, 1987,1988), Lambert et al. (1986), Olson and Richter (1979), Querci and Querci (1970), Smith andLambert (1990), Wannier et al. (1991)a Relative to solar composition.

b Assuming values of standard mean ocean water are representative of the solar isotopiccomposition.c Isotopic changes from solar that may be observable in presolar grains.

During the RGB stage, lasting about 500Ma for a solar mass star, the outer envelope becomes convective and penetrates into the region where partial H-

burning took place. Thereby, the unprocessed envelope material becomes polluted with the byproducts of the CNO-cycle (the “ first dredge-up ”).Observed 12C/13C ratios in RGB stars range from 6 to 20 (Table 8) - significantlyreduced from the (solar) main-sequence ratio of 89 - and indicate dredge-up of CNO-processed material. The N-isotopes are difficult to measure in stars.

31



Measured 17O/16O ratios in RGB stars are larger than solar, and 18O/16O ratios are below or near the solar ratio (Table 8).

Compared to standard evolution models, stars less massive than ~2.3 M ? show lower 12C/13C and Na and Al abundances than expected from the firstdredge-up of CNO-processed material (e.g., Lambert 1981, Gilroy and Brown1991, Sneden 1991). In order to bring nucleosynthesis models into accord withstellar observations, an extra mixing process was postulated (Charbonnel 1994,1995). This is known as cool bottom processing , “CBP” (Wasserburg et al.1995, Boothroyd and Sackmann 1999). In CBP, the base of the inert convectiveenvelope cycles through the regions that are on top of the H-burning shell so thatCN-processing of envelope material can occur. This leads to 12C/13C = 3.5 andhigh 14 N/ 15 N ratios in the observable stellar envelope. CBP is also assumed tooperate at later stages of stellar evolution (Nollett et al. 2003).

The RGB phase terminates when the He-rich core, left behind after main-sequence burning, ignites. This violent explosion is known as “He-core flash” for

stars <2.5 M ? . The He-core burning may occur concurrentlywith H-shell burning in stars of >2.5 M ? . The interior structureturns to a He-burning core, anoverlying H-burning shell, and aconvective envelope. At thisstage, lasting ~50 Ma in a 1 M ? star, the luminosity drops and theeffective temperature slightlyincreases so that the star “moves”from the RGB back to near themain sequence in the HR diagram. Helium-burning in thecore creates 12C by fusion of three 4He nuclei, known as the‘triple- α’ reaction: 4He ( 4He, γ)8Be( 4He, γ) 12C. In the mostmassive giant stars (~8-9 M ? )higher core temperatures maylead to α -capture on 12C to

produce 16O (i.e., 12C(4He,γ)16O).This creates a C- and O-rich ‘COcore’.Fig. 18. A cartoon of the different interior zones in

giant stars. The diagrams are not to scale and thezones where nucleosynthesis takes place are muchsmaller.

After He is exhausted in thecore, the star approaches theasymptotic giant branch (AGB).

32



This name comes from the location of these stars in the HR diagram, where AGBstars plot asymptotically to the red giant branch. Hydrogen-shell burning stops,and He-shell burning sets in. In larger mass stars ( >4-5 M ? ) the convectiveenvelope can penetrate deep enough to transport more CNO-processed material(rich in 4He, 13C, 14 N, 17O) previously made in the H-burning shell to the stellar surface (the ‘ second dredge-up’ ).

On the AGB, burning alternates between a thin H- and a He-shell surroundingthe inert core. The two shells are separated by an intershell region. If the He-

burning shell is ignited, temperatures raise and overlying zones expand.Expansion of the He-shell moves the intershell and the H-burning shell intocooler regions. Then H-burning is extinguished, the temperature-dependentenergy production in the expanding He-shell drops. At some point, He-burningceases, and subsequent contraction leads to ignition of a new H-burning shell.The successive re-ignitions of He-shell burning occur in intervals of ~10 4 a andare known as thermal pulses (TP).

Thermally pulsing ‘TP-AGB’ stars consist of an inert C-O-rich core,surrounded by a He-burning shell, a He-intershell layer, a H-burning shell, andthe convective envelope (Fig. 18). When H-shell burning is inactive, convectiontransports the carbon made by partial triple α-burning in the He-shell to thestellar surface (the ‘ third dredge-up’ ). This increases the observable C/O and12C/13C ratios. If a star experienced several thermal pulses and third-dredge-upepisodes it will show relatively larger amounts of He-burning products at itsobservable atmosphere.

The third dredge-up also brings up products made during H-shell burning,including 23 Na (the only stable Na isotope) and radioactive 26Al. These are madein the NeNa- and MgAl-cycles that are linked to the CNO-cycle (Forestini et al.1991, Gallino et al. 1994, Wasserburg et al. 1994, Guélin et al. 1995). Theobservation of 26Al in giant stars would provide another strong link to presolar grains because many different grains show evidence of incorporated 26Al.However, searches for 26Al in giant stars (using the isotopic shifts from 26Al and27Al in spectral bands of AlH, AlCl, or AlF, e.g., Branch and Perry 1970) give nofirm results. Guelin et al. (1995) obtained 26Al/ 27Al <0.04 for the closest, well-studied C-star CW Leo (IRC+10°216), but more measurements are required tocheck for 26Al in other giant stars.

Cool bottom processing may operate in low mass stars on the AGB during H-shell burning. This process is invoked to explain depletions in 18O/16O (relative tosolar) in some presolar grains (e.g., Wasserburg et al. 1995, Nittler et al. 1997),and it may also account for the production of 26Al (e.g., Nollett et al. 2003).

The TP-AGB stars show nucleosynthesis products made by neutron capturereactions, which are of particular interest here. The neutrons necessary for thisare produced by α- capture on 13C by the reaction 13C(α ,n) 16O. During a thirddredge-up episode, the H-rich envelope is in contact with the 12C- and He-rich

33



intershell. There 12C becomes abundant from the previous convective thermal pulse as fresh 12C is produced by partial triple alpha burning at the bottom of thethermal pulse and then spread into the He intershell. A small amount of protonslikely diffuses from the H-shell into the intershell, and with protons and 12C, the13C is made via by the reaction 12C(p, γ)13 N(β+ν)13C forming a “ 13C-pocket” at thetop of the intershell.

In intermediate-mass stars (>5 M ? ) the reaction 22 Ne( α ,n) 25Mg becomes amore important source of neutrons during later thermal pulses, and the neutrondensity strongly depends on the maximum bottom temperature in the convectiveshell. The 22 Ne for this neutron source comes from the reaction14 N( α,γ )18F(β+ν)18O(α,γ )22 Ne, which can occur in both low- and intermediate-mass stars. This reaction is the source of the Ne-E(H), which was one of the

beacons during the search for presolar grains. Lewis et al. (1990, 1994) argue that Ne-E(H) in SiC is unlikely to originate from decay of 22 Na (which only comesfrom novae or supernovae), because 22 Ne correlates with 4He, which is abundantin the He-rich shell.

Neutron capture by elements around the iron peak builds up the abundancesof heavy elements (e.g., Iben and Renzini 1983, Gallino et al. 1990, 1997, 1998,Busso et al. 1999). In AGB stars, neutron capture proceeds on a slow time scale(relative to the time scales of beta-decay of the radioactive nuclides produced)and is called the “ s-process”. Elements with relatively abundant stable isotopesmade by the s- process include Sr, Ba, Zr, Y, and the light rare earth elements, aswell as several isotopes of Kr and Xe. AGB stars that experienced many thirddredge-up episodes should show higher C/O and higher s-process elementabundances than stars like the sun or RGB stars of solar metallicity.

Another mechanism proposed to operate in stars of 4–7 M ? on the AGB ishot bottom burning (HBB) or envelope burning (Becker and Iben 1980, Renziniand Voli 1981, Boothroyd et al. 1995, Lattanzio and Forestini 1999). In suchmore massive stars, the base of the convective envelope is hotter and CNO-

processing can occur in the portion of the envelope located above the H-burningshell. This model predicts a decrease of 12C/13C to ~3.5 and increases of 14 N/ 15 Nup to 30000. The freshly produced 12C from triple α-burning in the He-shell isthus converted to 13C and 14 N. However, the C to N conversion may prevent anincrease in the C/O ratio and the AGB star cannot become a C-star, or dependingon timing, HBB could change a C-star back to an O-rich AGB star.

6.2. Types of cool giant stars

The chemistry of cool giant stars divides them into at least three major categories, with transitional types in-between. Starting with solar-like elementalcomposition and moving to more carbon and s-process element-rich objects, thegiant star spectral sequence is M-MS-S-SC-C. The M-stars, which comprise

34



two-thirds of giants, are either RGB stars (showing only products from the CNO-cycle) or AGB stars whose envelopes have just begun to become polluted by thethird dredge-up. The spectra of M-giants are dominated by metal oxide bandssuch as VO and TiO, which shows that these stars are relatively cool (~3000 K)near their surfaces. The S-stars are clearly on the AGB because they have C/Oratios near unity and over-abundances of s- process elements relative to solar. Inaddition to TiO and VO bands, S-star spectra show strong ZrO bands because Zr is increased by the s- process. The discovery of a pure s-process product,radioactive 99Tc (t 1/2= 2.1×10 5 a) by Merrill (1952) proves that the s-processindeed operates in these relatively low mass stars. Note that some S-stars areenriched in s-process elements but not in Tc (Little et al. 1987). Jorissen et al.(1993) call them “extrinsic S-stars” and use “intrinsic S-stars” for objects withTc. The extrinsic S-stars probably accreted the s-process elements from a moremassive companion star that already evolved through the AGB stage and is nowa white dwarf. Extrinsic S-stars may not be on the AGB.

Carbon stars atmospheres have C/O> 1 from multiple third dredge-upepisodes and they show bands from abundant carbon-bearing molecules such asC2, CH, and CN as well as atomic metal lines. Of the several sub-types of C-stars(N, J, R, and CH stars) only the N-type carbon stars are thought to be true TP-AGB stars. The evolutionary status of the other types of C-stars is not wellknown (Wallerstein and Knapp 1998).

Convective mixing brings the products of nucleosynthesis from the stellar interior to the stellar surface, but these products need to get into dust grains and

become part of the interstellar matter. This is accomplished by stellar winds thatdrive mass-loss from AGB star envelopes. The accumulation of gas andcondensates from the cooling gas creates circumstellar shells (Lafon andBerruyer 1991, Lamers and Cassinelli 1999). Sometimes “circumstellar shells”are referred to as “circumstellar envelopes” but here we use “circumstellar shells” to avoid confusion with the stellar envelope that directly surrounds theinterior nuclear burning zones. The shells can extend to several hundred stellar radii, with stellar radii themselves ranging from 100 to 500 R ? . They are dividedinto an inner shell in which thermochemical (e.g., condensation) processes aremore important, and an outer shell where UV-driven photochemical processesdominate (Glassgold 1996).

Circumstellar shells are detected by excesses in infrared radiation as lightfrom the central star is absorbed and re-emitted by dust at longer red and infraredwavelengths (e.g., Olnon et al. 1986, Gezari et al. 1993, Kwok et al. 1997). Insome instances, enshrouded stars are only detected by the infrared radiation fromtheir thick surrounding shells (e.g., Guglielmo et al. 1998, Volk et al. 2000).Typical mass-loss rates are ~10 -4 to 10 -7 M ? per year for M- and C-type AGBstars, and about a factor of ten less for S-stars (e.g., Loup et al. 1995, Olofsson etal. 1993). In contrast to the extensive mass-loss experienced by TP-AGB stars,

35



giants on the RGB do not show large infrared excesses, and their mass-loss is notvery productive (e.g., Jura 1999). Typical mass-loss rates of RGB stars are ~10 -9 to10 -8 M ? per year (e.g., Mauron and Guilain 1995). At some point, mass-losshas removed an AGB star’s envelope, and after several one to ten million yearsof alternating H- and He-shell burning in the inside and continuous mass-loss onthe outside, the AGB stage of stellar evolution comes to an end.

During the post-AGB stage (van Winckel 2003), mass-loss rates drop becausemost of the envelope is already gone. The circumstellar shell continues to expandwith terminal velocities of 10-30 km/s. As it separates more from the central star,the veil of dust and molecular absorptions around the hot stellar remnant is lifted.The last envelope matter leaves the star in a fast wind of 1000-4000 km/s andonce this much faster wind collides with the older, slower expanding material astrong shock front is created. The shock interaction of the ejecta and the radiationfrom the hot central star creates a planetary nebula, which continues to expandwith a final velocity of ~40 km/s (Kwok et al. 1978, Kwok 1994). One mayspeculate if the dust in the older shell is destroyed when it is hit by the shock wave. However, dust is still observed in planetary nebulae.

The remains of a giant star are a white dwarf consisting of the C-O-rich core,and a planetary nebula created from material of the former stellar envelope. Of aninitially 1 M ? star, about half remains, whereas stars of initially 4-8 M ? loosearound 80% of their main sequence mass to the ISM (e.g., Weidemann 2000).The local mass return from M-giants is estimated 1-2×10 -4 M ? /year/kpc 2 and C-giants provide about the same amount. The total mass returned to the galacticISM is ~0.3-0.6 M ? /year (e.g., Wallerstein and Knapp 1998).

6.3. Dust condensation in circumstellar shells

Dust grains condense in circumstellar shells around giant stars and then enter the ISM. Temperatures in the expanding circumstellar shell drop by adiabaticcooling, and, at a far enough distance from the star, become low enough (<2000K) for grain condensation. An important property of giant stars related to grainformation is stellar variability . Many AGB stars are “pulsating variable stars”and show more or less regular periodic variations in brightness (e.g., Hoffmeister et al. 1985). The variations are caused by thermal ionization and de-ionization

processes in the stellar envelope relatively close to the photosphere (e.g., Bowen1988). These processes regulate the opacity and energy transfer through theenvelope, and lead to radius expansions (dimmer star) and contractions (brighter star). The variability cycles last 100 to 600 d, with Mira-type variables showing

the most regular periods of 400-500 d (e.g., Hoffmeister et al. 1985). (Note thatstellar variability is not related to the thermal pulses at the TP-AGB stage, wherecyclic ignitions of He-shell burning occur at a frequency of about 10 4 a.)

36



Hence the dynamic nature of the envelope responsible for giant star variabilityalso affects the radial temperature and pressure structure within the circumstellar shell. Because condensation is a function of temperature and total pressure,stellar variability influences the absolute distance from the central star at whichgrain formation occurs. Fig. 19 shows how condensation may proceed during thevariability cycle (Lodders and Fegley 1997).

During maximum light, the star has the smallest radius and the highest photospheric temperature, and high temperatures within a few stellar radii do notfavor condensation. When the star changes phase, temperatures drop in theexpanding circumstellar shell. Dust formation should be most efficient at the lowtemperatures during minimum brightness. Because of the large radius expansion,the absolute grain-forming location is also further away from the center of thestar (but still within a few radii of the expanded star). After passing through theminimum phase, temperatures increase, the stellar radius decreases, and gas anddust from the circumstellar shell can “fall back” onto the star, which may lead todust evaporation. However, this can be prevented if dust grains are accelerated inthe circumstellar shell by radiation pressure from the star and leave into the ISM.

Considering the observedvariability periods (e.g., <600days for minimum to minimum),grain formation must occur ontime scales of at least half thevariability period (betweenminimum and maximum).Thenature of the dust condensingfrom ejected envelope materialdetermined by the C/O ratio inthe gas, which ranges from ~ 0.5(solar) to ~1.2 in giant stars.Changes in other elementabundances introduced bynucleosynthesis (e.g., higher Nabundances in RGB stars, andhigher s-process elementabundances in AGB stars) shouldnot affect the oxidation state.Under “oxidizing” conditions(C/O < 1) major element oxidesand silicates condense (e.g., Lord1965, Larimer 1969, Grossman1972, Ebel and Grossman 2000,Lodders 2003), and carbonaceous

Fig. 19. A sketch of grain condensation incircumstellar shells of giant stars during thevariability cycle.

37



dust forms under “reducing” conditions (C/O> 1) from a gas of otherwise solar elemental composition (Friedemann and Schmidt 1967, Friedemann 1969,Gilman 1969, Larimer 1975, Lattimer et al. 1978, Larimer and Bartholomay1979, Lodders and Fegley 1995, 1997).

Table 9 lists condensates of the more abundant elements in circumstellar shells of M-and C-stars expected from thermochemistry, the minerals seen incircumstellar shells and planetary nebulae, and the known presolar minerals.Several minerals are both known in circumstellar environments and as presolar grains and are the first condensates to appear under oxidizing and reducingconditions, respectively. The condensation temperatures as a function of total

pressure are shown in Figs. 20 and 21. The range in total pressures shown coversthat expected in circumstellar shells, e.g., 10 -5 to 10 -8 bar at T<2000 K. With thenotable exception of carbon, the condensation temperatures generally increasewith increasing total pressure.

Condensates for M-stars are similar as for a solar composition gas becausetheir C/O ratios result in comparable oxidation states (e.g. Dominy et al. 1986,Smith and Lambert 1985, 1986, 1990, Smith et al. 1987). Condensates of Ca andAl include corundum, spinel, hibonite, gehlenite, and anorthite. Silicon and Mgcondense as forsterite and enstatite, and Fe forms an FeNi-alloy and troilite.

Several of these condensates are present in circumstellar shellsand in planetary nebulae byinfrared and mid-infraredobservations. Abundant silicategrains are expected in O-richshells, and the 9.7 µ m emissionin M-giant spectra observed byGillett et al. (1968) was ascribedto silicates by Woolf and Ney(1969). This silicate feature isobserved in shells of many O-rich stars (e.g., Little-Mareninand Little 1988, 1990, Sloan etal. 1996). The 20 and 28 µ memissions identify Mg-rich

pyroxene and forsteritic olivine(e.g., Molster et al. 2002a-c, Suh2002). Circumstellar silicatesseem to be both crystalline andamorphous, with thicker andcooler shells favoring crystallinesilicates.

Fig. 20. Condensation temperatures as a function of total pressure for a gas with C/O = 0.5.

38



Schmid-Burgk and Scholz (1981) already discussed Al 2O3 formation in M-giants and an emission feature at 13 µ m, present in 40-50% of all O-rich AGBstars (Sloan et al. 1996), was identified as corundum (Onaka et al. 1989,Begemann et al. 1997, Kozasa and Sogawa 1997, Fabian et al. 2001). Corundummay have an additional emission at 21 µ m (Begemann et al. (1997), which isseen in several post-AGB objects. However, in these stars the 21 µ m feature isnot accompanied by one at 13 µ m, and the origin of the 21 µ m emission remainsmysterious. Another possible source for the 13 µ m feature is magnesian spinel,especially if it is accompanied by 16.8 and 32 µ m emissions (e.g., Posch et al.1999, Fabian et al. 2001). However, the identification of spinel is uncertain

because silicate emissions may interfere and other interpretations were suggested(e.g., Sloan et al. 2003).

Metallic iron, an expected major dust component, is infrared inactive. Kemper et al. (2002) infer its presence in the dusty circumstellar shell of an evolved O-rich giant because an abundant infrared opacity source is needed to fit the

spectrum. Troilite, a lowtemperature condensate, is not yetidentified in M-giants. Perovskiteis also not yet detectedcircumstellar shells. Posch et al.(2003) found that the infraredsignatures of perovskite, other Ca-titanates, and Ti-oxides coincidewith those from silicates andaluminum oxides. More abundantSi and Al condensates likely reducethe spectral band strengths of Ti-

bearing condensates so that Ti- bearing dust may not be resolvable.

Condensates for reducingconditions (Fig. 21) are calculatedfor the average composition of N-stars by Lambert et al. (1986).Condensates are carbon andcarbides (SiC, TiC), sulfides (CaS,MgS), and nitrides (AlN, TiN).Iron condenses as silicide (FeSi)and troilite. Silicates also condense

at C/O >1, despite a persistent misconception among some researchers thatsilicates do not form under reducing conditions. (The reduced enstatite chondriteswith mainly pure enstatite show that Mg-silicates formed at C/O >1, e.g., Krot etal. 2000).

Fig. 21. Condensation temperatures as a functionof total pressure at C/O = 1.1.

39



The 11.3 µ m emission of SiC is prominent in circumstellar shells of C-stars(e.g., Gilra 1971, Treffers and Cohen 1974, Forrest et al. 1975, Little-Marenin1986, Little-Marenin et al. 1987). Carbon, either amorphous and/or in graphiticform, is believed to be present around many C-stars (Rowan-Robinson and Harris1983, Jura 1986, Martin and Rogers 1987, Orofino et al. 1987, Blanco et al.1994). Diamond is not an expected condensate under equilibrium thermodynamicconditions. However, emissions of surface-hydrogenated diamond at 3.43 and3.53 µ m (Guillois et al. 1999) are seen in the C-rich shell of HR4049, a verymetal-poor post-AGB star (Geballe et al. 1989, van Kerckhoven et al. 2002), andin the proto-planetary nebula CRL 2688 (Geballe et al. 1992).

The 30 µ m emission observed in C-rich AGB stars and planetary nebulae isattributed to MgS (Goebel and Moseley 1985, Nuth et al. 1985, Hony et al.2002a, Hony and Bouwman 2004). Another expected abundant condensate, FeS,was identified in planetary nebulae (Hony et al. 2002b, Begemann et al. 1994).

The phase responsible for the 21 µ m emission from several post-AGB starsand planetary nebulae mentioned before has been a puzzle ever since it was firstdescribed (Kwok et al. 1989). Among other suggested phases (e.g., spinel,nanodiamonds, SiS 2, polyaromatic hydrocarbons), this feature was ascribed toTiC nano-particles (von Helden et al. 2000), but this assignment is debated (e.g.,Li 2003). Although TiC is one of the first condensates from a reducing gas, thelow Ti abundance (relative to Si or C) may prevent detection of TiC incircumstellar shells. In addition, presolar graphite grains contain TiC asinclusions, which implies that detection of TiC in circumstellar shells is limited

because TiC is hidden within other grains.The total pressure and temperature conditions under which presolar grains

formed can be estimated from thermochemical calculations. Many transitionelement carbides (e.g., TiC, ZrC, MoC) are quite refractory (e.g., Larimer 1975,Lattimer et al. 1978, Lodders and Fegley 1995, 1997) and tiny inclusions of suchcarbides within graphite grains apparently acted as nucleation seeds for thesurrounding carbon (Bernatowicz et al. 1991, 1996, Croat et al. 2003). The

presence of refractory carbide inclusions in graphite and the absence of SiCinclusion corresponds to the condensation sequence transition metal carbides-C-SiC. A detailed study using equilibrium and non-equilibrium chemistry byBernatowicz et al. (1996) shows that favorable conditions for the observedcondensation sequence of metal carbides and carbon in the stellar outflowsrequire total pressures >10 -7 bar and C/O > 1.05. Chigai et al. (1999) considerednon-equilibrium conditions and inferred C/O ratios less than 1.26 to 1.46 andtotal pressures between 10 -7 to 2×10 -9 bars.

40



Table 9. Expected and observed major element condensates for giant stars and presolar grains a M-stars C-stars

Abundance b Mineral and idealformula

Stellar shell

Presolar Grains

Mineral and idealformula

Stellar shell

PresolarGrains

O 1.41×10 7 oxides and silicates silicatesC 7.08×10 6 — titanium carbide TiC

graphite Csilicon carbide SiCdiamond?

√√√ √(?)

√√√ √

N 1.95×106 — osbornite TiNaluminum nitride AlN

Mg 1.02×10 6 spinel MgAl 2O4 forsterite Mg 2 SiO 4 enstatite MgSiO 3

√ √ √

√ √ √

niningerite MgSspinel MgAl 2O4 forsterite Mg 2SiO 4 enstatite MgSiO 3

√√√√

√√ √

Si 1.00×10 6 gehlenite Ca 2Al2SiO 7 forsterite Mg 2SiO 4 enstatite MgSiO 3

√ √

√ √

silicon carbide SiCiron silicide FeSiforsterite Mg 2SiO 4 enstatite MgSiO 3

√

√ √

√

√ √

Fe 8.38×10 5 iron alloy FeNischreibersite (Fe,Ni) 3P

iron silicide FeSiiron alloy FeNi

S 4.45×10 5 troilite FeS oldhamite CaSniningerite MgStroilite FeS

√ √

Al 8.41×10 4 corundum Al 2O3 hibonite CaAl 12O19 grossite CaAl 4O7 gehlenite Ca 2Al2SiO 7 spinel MgAl 2O4 anorthite CaAl 2Si2O8

√

√

√ √

√

aluminum nitride AlNcorundum Al 2O3 spinel MgAl 2O4 anorthite CaAl 2Si2O8

√ √

√ √

Ca 6.29×10 4 hibonite CaAl 12O19 grossite CaAl

4O

7

gehlenite Ca 2Al2SiO 7 anorthite CaAl 2Si2O8

√ oldhamite CaSanorthite CaAl

2Si

2O

8

Na 5.75×104 albite NaAlSi 3O8 albite NaAlSi 3O8

halite NaCl (?) Ni 4.78×104 kamacite & taenite

schreibersite (Fe,Ni) 3P√ kamacite & taenite

schreibersite (Fe,Ni) 3P√

Cr 1.29×10 4 Cr in FeNi alloy daubréelite FeCr 2S4 Mn 9.17×10 3 Mn 2SiO 4 in olivine alabandite (Mn,Fe)SP 8.37×10 3 schreibersite (Fe,Ni) 3P schreibersite (Fe,Ni) 3PCl 5.24×10 3 sodalite Na 4[AlSiO 4]3Cl halite NaCl (?)K 3.69×10 3 orthoclase KAlSi 3O8 orthoclase KAlSi 3O8

sylvite KCl (?)Ti 2.42×10 3 perovskite CaTiO 3 titanium carbide TiC

osbornite TiN√ (?) √

a Adapted from Lodders and Fegley (1999). Elements are in order of decreasing solar abundance.Condensates are in order of appearance with decreasing temperature for each element. A “?” meansthat entry is uncertain. b solar photospheric abundances where Si = 1×10 6 atoms (Lodders 2003).

41



Overall, thermodynamic equilibrium calculations are very useful for modelingthe chemical composition of dust from stars. We also refer the reader to thereview by Sedlmayr and Krüger (1997) for non-equilibrium and kineticconsiderations.

6.4. Synthesis: Presolar grains from RGB and AGB stars

The known presolar minerals from giant stars are as follows. Corundum,spinel, and silicate grains of the oxide groups I and III come from RGB stars.AGB stars contributed corundum, spinel, and hibonite grains to the oxide groupsI, II and III; mainstream, Y, and Z SiC grains; probably graphite grains with12C/13C ratios larger than solar, which comprise more than 70% of the high-density KFC1 graphite fraction; and possibly some diamonds.

6.4.1. Grains from red giant branch (RGB) stars

Dust from RGB stars should consist of oxides and silicates. After the firstdredge-up, diagnostic changes in isotopic compositions relative to solar areexpected in C, N, and O from processing by the CNO-cycle. Oxidizedcondensates should not contain much C and N (possibly as impurity, if any), soonly the O-isotopes are useful to sort out the presolar corundum, spinel, hibonite,and silicate grains from RGB stars. After the first (and second) dredge-up,17O/16O should increase over the solar ratio, but the 18O/16O ratio should remainnear the solar ratio in M stars on the RGB and early AGB (Table 8). Thisindicates that grains from RGB stars are to be found in group I and III in Fig. 14.However, oxide grains from RGB and early AGB stars should not have evidenceof 26Al. The oxide grains measured for Al-Mg systematics (Fig. 15) show that26

Al is absent in 64% of the grains from group III and in 30% of the grains fromgroup I. Apparently group III is dominated by oxides from RGB stars .

6.4.2. Presolar grains from asymptotic giant branch (AGB) stars

A star on the AGB produces either oxidized or reduced condensates,depending upon how many third dredge-up episodes have occurred to change theC/O ratio from the near-solar RGB ratio to ratios above unity. Major changes onthe AGB with respect to the RGB stage are increases in C and s-process elementabundances and rising 12C/13C ratios from the RGB values below ~20 (e.g., Smithand Lambert 1985, 1986). Nitrogen-rich in 14 N is mainly inherited from the RGBstage. Another difference between TP-AGB and RGB (and early AGB) stars isthat 26Al made during H-shell burning is brought to the surface during the thirddredge-up episodes. Neutron capture reactions in AGB stars also modify the Caand Ti isotopes, and increases in the isotopes 47Ti, 49Ti, and 50Ti relative to 48Tiand 42Ca, 43Ca, 44Ca, and 46Ca relative to 40Ca are expected (e.g., Hoppe et al.1994, Choi et al. 1998, Amari et al. 2000a) but presolar oxides from RGB starsshould contain the unaltered initial Ca and Ti isotopes of the stellar source.

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Dust from O-rich AGB stars includes the expected condensates corundumand spinel, however, the presolar oxide grains of groups I and III plotted in theO-isotope diagram (Figs. 14, 15) are a mix of grains from RGB and AGB stars.The inferred presence of 26Al shows that group I has more grains from AGB stars(70%) than group III (36%). However, these percentages do not reflect therelative distribution of RGB and ABG stars contributing to presolar oxides

because one AGB star produces much more dust than one RGB star.Measured O-isotope ratios in RGB and AGB stars and planetary nebulae are

compared to oxide grain data in Fig. 14. RGB stars have about normal 18O/16Oand higher 17O/16O and plot together with group I grains. Data for AGB stars onlyinclude stars that show Tc, which ensures that they are truly on the AGB. SeveralAGB stars and PN plot with group I grains; however, there is a notable shift tolower 18O/16O for AGB stars (but not PN) towards group II grains.

The origin of oxide grains in groups II and III (Figs. 14, 15) was also ascribedto AGB stars (Nittler et al. 1997). Group II grains are slightly enriched in 17O andheavily depleted in 18O relative to solar, and many of these grains have thehighest inferred 26Al content among oxide grains. The depletion of 18O and thelarge 26Al abundances in group II oxide grains is plausibly due to cool bottom

processing in intermediate mass giants (<3M ☼ ). However, there is no detection of 26Al in O-rich AGB stars and it is odd that the O-isotopic compositions in AGBstars do not overlap more with the field spanned by the group II grains (Figs. 14,15). None of the stellar O-isotopic compositions overlap with grains of group IIIand IV, and the overlap is marginal at best for group II grains. This lack of overlap with data for current AGB stars could favor the interpretation that grainsof group II and III reflect chemical galactic evolution processes (Nittler et al.1997).

The O-isotopic compositions of the two hibonite grains found by Choi et al.(1999) favor an origin from a low mass (~1.7M ? ) AGB star of solar metallicityand a ~1.2M ? RGB or early AGB star. The inferred initial 26Al/ 27Al ratiosindicate that H-shell burning and CBP took place in the grains’ stellar sources.

Dust from C-rich ABG stars includes most SiC grains, some yet unkownfraction of the graphite grains, and possibly – if any - some of the nanodiamonds.

A little more than 90% of all presolar SiC originated from AGB stars, mainlythe mainstream SiC and related Y and Z grains (e.g., Hoppe et al. 1994, 1997,Hoppe and Ott 1997, Amari et al. 2001b). Their chemical and isotopiccompositions are in accord with abundances in N-stars and AGB models. Withinuncertainties, there is good overlap in the C and N isotopic compositions of mainstream SiC grains, N-stars and planetary nebulae (Fig. 6). The J-type C-starsoverlap with data of the A+B grains (see below).

The paucity of Si-isotopic determinations in stars (Fig. 7) makes a comparisonto SiC grains difficult. The two determinations for the C-star IRC+10°216 fallinto the array of SiC mainstream grains. The other data are for red giants not yet

43



on the AGB. Unfortunately, these measurements are quite uncertain. Leaving thisaside, the lack of overlap between the Si-isotopic composition of SiC grains andRGB stars is expected because the ejecta from RGB stars are not reduced enoughfor SiC condensation. In this context we note that future measurements of Siisotopes in presolar silicates may show some overlap if these originate in RGBstars. However, models indicate that the Si-isotopic composition is not affectedduring nucleosynthesis in RGB stars and that it is probably only mildly modifiedin AGB stars of solar metallicity. Then the observed Si-isotopic composition isthat with which a star was born and observed differences in Si-isotopiccompositions point to galactic chemical evolution (Zinner et al. 2001).

The distribution in 12C/13C ratios of SiC (Fig. 22) corresponds well to that of C-stars determined by Lambert et al. (1986). The 12C/13C distributions measuredin C-stars by Ohnaka and Tsuji (1996, 1999) and Schöier and Olofsson (2000)are similar but their maxima are shifted to lower 12C/13C ratios. However, thedata by Lambert et al. (1986) are regarded as more reliable. The similarity in the12C/13C distribution of SiC grains with that of N-stars is consistent with theconclusion that most of the SiC grains (93%) are from evolved AGB stars.

The trace element abundances in mainstream SiC grains provide a close link to N-stars because these are also enriched in s-process elements. The differentabundance patterns for mainstream SiC grains (Fig. 9) and the fractionationsamong the elements can be understood by considering fractional condensation.Lodders and Fegley (1995) calculated the trace element abundance patternsobserved in SiC grains (Amari et al. 1995c) with the goal to find the elementalabundances in the stellar sources, which are necessary to explain the patterns byfractional condensation. This then constrains the types of C-stars that could have

produced the SiC grains. The patterns in Fig. 9 require source abundances of s- process elements ranging from solar up to 10×solar, and in some cases, theremoval of ultrarefractory trace element carbides prior to SiC condensation. TheY grains also show enrichments in s-process elements, and their abundance

patterns require condensation from a gas enriched in s-process element by afactor of 10 (Lodders and Fegley 1995). Such a large enrichment is consistentwith the suggested origin of Y grains from low-metallicity AGB stars, wherehigher s-process yields are expected (Amari et al. 2001b).

Several of the noble gas isotope anomalies (e.g., the AGB “G” componentsfor the noble gases in Fig. 5) can be accounted for if SiC grains formed aroundAGB stars. Early on, the G component of Kr was found to be strikingly similar incomposition to theoretical expectations from s-process nucleosynthesis (Gallinoet al. 1990, 1997). However, the Kr isotopes still pose some challenges tonucleosynthesis and stellar models. For example, why do the abundances of nucleosynthetic products, which are determined by the conditions in the hotinterior of AGB stars, correlate with grain size, which reflects the conditions inthe cool, expanding circumstellar shell outside? Lewis et al. (1990, 1994)

44



observed a positive correlation of the s-process 86Kr/ 84Kr ratio of theG component with increasingaverage grain size for the differentSiC aggregates. In contrast, the80Kr/ 84Kr ratios decrease withaverage aggregate grain size.Verchovsky et al. (2004) proposedthat the grain size correlation for Kr isotope abundance ratios can beexplained by assuming that onelow-energy Kr component wasimplanted into the grains during theAGB stage and a second, high-energy component was incorporatedat the post-AGB stage during

planetary nebula formation, whichwas too energetic to be captured bysmall grains.

The larger 86Kr/ 84Kr ratios in thelarger-size grains fractions requiredlarger neutron exposures for the s-

process. Theoretical models(Gallino et al. 1990, 1997) suggestthat higher neutron exposures areachieved in low-metallicity AGBstars, which then may suggests thatthe larger grains formed in low-metallicity stars at higher neutronexposure than the finer grains.However, this does not yet explainwhy low-metallicity stars producelarger grains in their outflows, andthe effects of stellar mass andmetallicity on grain sizedistributions are yet to be explored.

There is also a correlation between the Sr and Ba isotopesmade by the s-process and average SiC grain size. However, the trend is opposite to that observed for s-process Kr: The 88Sr/ 86Sr and 138Ba/ 136Ba ratios decreasewith increasing average aggregate grain size, which implies that neutronexposure decreased with grain size. A probable explanation is that grains

Fig. 22. Distribution of 12C/13C ratios in presolar SiC (Hoppe et al. 1994), presolar graphite (Hoppeet al. 1995, Travaglio et al. 1999, Amari,unpublished), and in C-stars (Lambert et al.1986). Graphite and SiC data extend to 12C/13C of several thousand (Figs. 6, 8). Only the observedrange for C-stars is covered here.

45



carrying the s-pro cess Kr are not the same grains that carry the s-process Sr and/or Ba. Nichols et al. (2005) analyzed He and Ne in single SiC grains andfound only four percent of the grains are gas-rich. Therefore, it is notunreasonable to assume that Kr-rich grains may be in a different population thanthe Sr- and Ba enriched grains. However, more correlated measurements for Sr,Ba, and noble gases in individual grains are required to sort out all theseobservations.

The link of mainstream SiC grains to carbon-rich AGB stars also keepsgrowing from isotopic measurements of heavy elements, such as Mo, Zr, and Ba,in individual grains (Nicolussi et al. 1997, 1998a, Savina et al. 2003b). Oneinteresting recent result is the measurement of Ru isotopes which indicate thatTc, the herald of the s-process in AGB stars, condensed into presolar mainstreamSiC grains (Savina et al. 2004).

The identification of presolar graphite from AGB stars is more difficultthan for presolar SiC. The excellent agreement in the C-isotopic distributions(Fig. 22) among C-stars and SiC grains (mostly mainstream grains) stronglysupports that mainstream grains formed around AGB stars very similar to most

N-type C-stars we presently observe. If there is a major contribution of graphitegrains from N-stars, we expect that the distributions for graphite grains show a“peak” that coincides with that of N-stars, as is seen for SiC grains. However, thedistributions for graphites from the density fractions KFA1, KFB1, and KFC1 donot correlate with that of present-day N-stars. On the other hand, many graphitegrains in the KFC1 fraction may have originated from low-metallicity AGB starsfor which another 12C/13C distribution is expected. The high inferred 86Kr/ 83Kr =4.80 (Amari et al. 1995a) and the high 12C/13C ratios (a few hundred or evenabove a thousand, Fig. 11) in many KFC1 graphites are consistent with theexpected values for low-metallicity AGB stars (Gallino, private communication).In addition, many graphite grains from the high-density separate KFC1 containrefractory carbides with compositions ranging from nearly pure TiC to nearly

pure Zr-Mo carbides (Bernatowicz et al. 1996). The abundances of Zr and Morelative to Ti are larger than the solar ratio and indicate enrichments of the s-

process elements Zr and Mo in the stellar sources. However, it is currentlydifficult to pin down the fractional abundance of “AGB” graphite grains in thehigh-density fraction.

In Fig. 22, the distributions for graphites from the density fractions KFA1,KFB1, and KFC1 show a spike near 12C/13C = 89, the solar value. Such graphitegrains could have a solar system origin or an origin in sources that coincidedwith, or better, dominated the solar system’s C-isotopic composition. The KFB1and KFC1 fractions show another major peak for 12C/13C < 20, similar to A+BSiC grains and J-stars. Only the low-density fraction KE3 has a distribution thatmay correspond to N-stars (Fig. 22). However, Si- and O-isotopes relate many of

46



the KE3 grains to supernovae. This illustrates that more than one criterion isrequired to clearly relate presolar grains to their parent stars.

The presence of diamonds from C-rich AGB stars among presolar diamonds cannot be excluded but currently also cannot be firmly concluded. If the spectroscopic assignments to diamond are correct, nanodiamonds seem to be

present in some post-AGB stars and planetary nebulae. In that case, the knownSiC and graphite from AGB stars should be accompanied by AGBnanodiamonds. However, more chemical and isotopic information on presolar nanodiamonds is required to establish a link to AGB sources.

6.5. Presolar grains with signatures akin to AGB stars but an imperfectmatch

The origin of the second largest SiC population, the A+B grains, remainsenigmatic (Amari et al. 2001a). However, some types of C-stars must beresponsible for their production. More than one type of stellar source is necessaryto explain the spread in 14 N/ 15 N ratios (40–10 4) and the two types of A+B grainsdistinguished by their trace element concentrations. The A+B grains have either abundance patterns consistent with relative solar abundances of s-processelements (Fig. 9), or patterns like mainstream grains that require enhancement of

s-process elements at the stellar sources (Amari et al. 1995c, 2001a, Lodders andFegley 1995, 1998).

A comparison of the C- and N-isotopic composition of A+B grains and C-stars (Fig. 6) shows that J-type carbon stars have similar isotopic compositions.The C-isotope histograms (Fig. 22) also point in this direction, and the lack of s-

process element enrichments (Utsumi 1970, 1985, Kilston 1975, Abia and Isern

2000) makes J-stars good candidates for the A+B grains.The J-stars have C, N, and O elemental abundances similar to N-stars, but

lower 12C/13C isotopic ratios (Lambert et al. 1986, Abia and Isern 1997, Ohnakaand Tsuji 1999). The low 12C/13C ratio and the absence of s-process enrichmentsin J-stars is a long- standing mystery, which now extends to the A+B grains. Theevolutionary status of J-stars is unclear and they are unlikely to be on the AGB(e.g., Lloyd-Evans 1991, Abia and Isern 2000). According to nucleosynthesismodels, 12C, which is necessary to make a stellar atmosphere rich in carbon, is

produced during the AGB stage of low to intermediate mass stars. Without anymodifications, this also increases the 12C/13C ratio, but J-stars have the lowestvalues among C-stars. Furthermore, the s-process operates during the AGB stage,

but there are no s-process element enrichments in the J-stars (e.g., Utsumi 1970,

1985, Abia and Isern 2000). The situation to explain the A+B grains with s- process enrichments is also far from being settled (Amari et al. 2001a). Raretypes of stars that can produce reduced dust are possible sources for these grains.These include low-metallicity CH stars that evolve in binary systems, and stars

47



such as Sakurai’s object, which appears to be a “born-again” AGB-like star dueto external mass accretion (Amari et al. 2001a, Lodders and Fegley 1997, 1998).

7. Massive stars and supernovae (SNe)

Astronomical observations of supernovae and their remnants reveal the presenceof dust, so they are potential sources of presolar grains. For a long time, searchesfor infrared excesses indicative of dust only showed small amounts of dust, butrecent sub-millimeter observations indicate large amounts of very cold dust insupernova remnants (Dunne et al. 2003). Supernovae now seem to be confirmeddust producers on a similar scale as giant stars (e.g., Dunne et al. 2003, Morganet al. 2003, Kemper et al. 2004) . One would suspect that SNe produce largeamounts of dust because they supply the major elements (e.g., Mg, Al, Si, Ca,

Fe) to the ISM and increase the galactic metal content with time. The major elements will eventually condense, so there must be dust from supernova ejecta,and, by implication, dust grains from SNe among presolar grains.

There are different types of supernovae. Ejecta of SNe type I do not showatomic H lines in their optical spectra, whereas SNe of type II do. This indicatesthat a star that turns into a SN type I must have become very H-poor relative tothe H-rich composition of normal main sequence and giant stars. Type I SNe fallinto several subtypes (e.g., Ia, Ib, Ic), and type Ia is special among all SNe

because it involves a binary system.

7.1. Evolution of massive stars to supernovae

The evolution of massive stars (> 8-9 M ? ) is described in detail by Maeder (1990), Woosley and Weaver (1995), Arnett (1996), Pagel (1997), Wallerstein etal. (1997), Rauscher et al. (2002), Woosley et al. (2002), Meynet and Maeder (2003) and Truran and Heger (2004). The initial mass of a star and its mass-losshistory determines how the star ends its life. The AGB stars end up as whitedwarfs because mass-loss does not leave enough mass (<1.4 M ? ) to sustainnuclear reactions in the C-O-rich core left from He-core burning. The H- and He-shell burning also stops because the envelope is lost. The evolution of moremassive stars (>8-9 M ? ) initially proceeds similar to that of low- andintermediate stars but occurs more quickly, and nucleosynthesis continues inseveral subsequent steps.

Before they go supernovae, massive stars develop a concentric shell structurein composition (Fig. 23) as a result from the succession of the major nuclear

burning stages (see reviews noted above). Each major burning stage leads toashes that provide the fuel for the next stage. For example, the first major reactions in the stellar core are H-burning to He, and then He-burning to 12C and

48



16O, which also happens in less massive AGB stars. After this, carbon burns inthe core:

12C ( 12C,α ) 20 Ne

followed by Ne-burning which starts by photodisintegration of 20 Ne andsubsequent α− capture by 20 Ne:

20 Ne ( γ,α )16O20 Ne( α,γ )24Mg( α,γ )28Si

followed by O-burning, which produces e.g., 28Si, 32S, 36Ar, 40Ca:16O(16O,α) 28Si

followed by Si-burning initiated by photodisintegration of 28Si:28

Si( γ,α )24

MgThis is followed by α -, neutron- and proton-capture reactions on 28Si, which build iron peak elements up to 56 Ni.

All core-burning stages are accompanied by burning in shells of a similar succession, e.g., H-shell burning and He-core burning, He-shell burning and C-core burning; C-shell burning and C-O core burning, etc. Compared to thetimescales of H- and He-core burning, the next core-burning stages are greatlyaccelerated: carbon-burning lasts for a few hundred years, and Si-burningaround a day. After Si is exhaustedas fuel in the core, the core collapses,and a resulting shockwave drives the

supernova explosion. At this stage,explosive nucleosynthesis takes place. In the innermost zone, nuclear statistical equilibrium producesradioactive 44Ti, 48Cr, 49V, 51Mn,52Fe, 55Co, 56 Ni, 59Cu, and 60Zn,which decay to 44Ca, 48Ti, 49Ti, 51V,52Cr, 55Mn, 56Fe, 59Co, and 60Fe,respectively. These isotopes are verycharacteristic of SNe. Aside from56Co and 57Co in SN 1987A, γ-linesfrom 44Ti are reported in the Cas ASN remnant (Iyudin et al. 1994), andobserved excesses of 44Ca (the decay

product of 44Ti) provide a direct link of some presolar grains tosupernovae.

Fig. 23. The compositional zones resulting fromnucleosynthesis in massive stars before they gosupernovae are conveniently visualized in an“onion shell” diagram. Zones are labeled with theelements that make the major nuclear fuel and

products. Some nuclear reactions and products of interest to presolar grain studies are indicated.

49



Heavy element nucleosynthesis by the r - and p-processes is generallyassociated with SN explosions, but locating the sites for these processes is still

problematic (e.g., Wallerstein et al. 1997, Woosley et al. 2002). The rapid

neutron capture process ( r -process) requires high neutron densities and hightemperatures so that neutron capture on seed nuclei can proceed on a faster timescale than beta decay of intermediately produced radioactive nuclei. The p-

process is responsible for building proton-rich nuclei that cannot be produced

Fig. 24. The isotopic ratios for C, N, O, Al, and Si within a 20 solar mass star directly beforethe supernova explosion (top) and after the passage of the supernova shockwave (bottom). Theshell structure (Fig. 23) is indicated at the top and jumps in the isotopic ratios clearly separatesome of the individual zones. The mass plotted on the x-axis gives the total mass of the star that lies under a given zone. The data are from calculations by Rauscher et al. (2002).

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from nuclei made by the s- or r -processes and subsequent beta decay. In thecontext here, the anomalies in the p- and r -process Xe isotopes (Xe-HL) weremost diagnostic of supernova contributions to presolar grains.

The compositions of pre-supernova stars and supernovae have been calculated by Woosley and Weaver (1995), Thielemann et al. (1996) and Rauscher et al.(2002). Fig. 24 shows the compositional “cross-section” through a 20 M ? star (Rauscher et al. 2002) for the major isotope ratios that are of interest for presolar grain studies. The individual zones are indicated at the top, and the zone“boundaries” are characterized by sharp changes in isotopic compositions. Theimportant issue here is that some different zones must mix to some extent duringthe supernova explosions in order to produce the observed isotopic ratios in

presolar grains. In particular, products from the He/N, He/C, and the innermostlayers must mix, but the mixing cannot be complete because the C/O ratio must

be sufficiently high to allow formation of reduced dust. The innermost zones areall rich in O so C/O < 1, and only the He/C and He/N zones contain enough C toachieve C/O > 1, which leads to SiC and graphite condensation.

Stars with masses around 30 - 40 M ? lose their H-rich zones in a slow windduring their red super-giant phase. Stars above ~40 M ? shed their H- and He-richzones in a fast wind during the main sequence stage, from where they evolve toluminous blue variables and then to Wolf-Rayet (WR) stars, and these verymassive stars may skip the red (super) giant phase (Maeder 1990, Meynet andMaeder 2003). The substantial mass-loss produces extensive gas and dust shellsaround WR stars and they are therefore candidates for producing presolar grains.

Once the H-rich zone is lost, the products from H-burning by the CNO-cycle become visible as a N-rich “WN” Wolf-Rayet star. In the more massive carbon-rich Wolf-Rayet (WC) stars, loss of the H- and He-rich zones opens the view tothe 12C-rich layer generated from triple α- burning. In O-rich WO stars (likelymore massive than WN and WC-stars), α-capture on carbon has increased the Oabundance. The circumstellar shells of WR stars may contain O-rich dust (inWN, WO), and reduced dust (in WC), depending on how much of the H- and He-zones were lost. Relative to solar, presolar grains from ejected envelope materialof Wolf-Rayet stars should show a few of the following characteristics:enrichments in 14 N and 22 Ne, evidence for 26Al from the CNO-, NeNa- andMgAl-cycles during the WN-stage, 41Ca from neutron capture, and 12Cenrichments from He-burning (e.g., Arnould et al. 1997).

7.2. Dust from supernovae

There is not much information yet about of the mineralogy of dust fromobservations of supernova ejecta and remnants. An emission feature at 22 µ mseen in two supernova remnants is interpreted as “Mg-protosilicate” (Arendt etal. 1999, Chan and Onaka 2000), and a small feature at 13 µ m may be attributed

51



to corundum. However, the expected dust composition can be modeled if elemental abundances for these dust producers are available. In absence of observational data, results from nucleosynthetic network computations andhydrodynamic codes (e.g., Woosley and Weaver 1995, Meyer et al. 1995,Thielemann et al. 1996, Rauscher et al. 2002) are good resources for elementalabundances in SNe of different masses and in individual zones.

There are three possibilities to model supernova condensates. First, one caninvestigate the types of condensates that are expected from each individual zonecomposition of a supernova. The second is to take an average, overall ejectacomposition to calculate the types of condensates. The third is to assume more or less selective mixing of matter between different zones. The reason for thesedifferent approaches is that the extent of mixing in supernova ejecta is not verywell known. Observations and hydrodynamic models (e.g., Ebisuzaki andShibazaki 1988, Herant et al. 1994, Hughes et al. 2000, Kifonidis et al. 2003)indicate that the ejecta are relatively well mixed. On the other hand, limitedmixing between the various zones of a supernova is required to explain theisotope data of presolar grains, if nucleosynthesis predictions for the zonecompositions are correct (e.g., Travaglio et al. 1999).

The condensates expected from the individual supernova shells and from theoverall ejecta composition was computed by Lattimer et al. (1978). Despite thewide range in individual zone compositions, the condensate chemistry in SNe isnot that much different than described above for giant stars and is mainlygoverned by the C/O ratio so that Table 9 also can be used as an approximateguide to the major expected supernova condensates.

Kozasa et al. (1989a,b, 1991) utilized the extensive observations that followedthe explosion of SN 1987A and compared them to their thermochemical andkinetic calculations. Depending on the degree of mixing of material fromdifferent zones, they inferred that graphite, corundum, enstatite, and magnetiteare expected in the ejecta. Ebel and Grossman (2001) address condensation insupernova ejecta but focus on the more specific question of whether or notreduced condensates can form at C/O < 1 if formation of CO is suppressed, assuggested by Clayton et al. (1999, 2001).

7.3. Synthesis: Presolar grains from supernovae

When searching for presolar grains from SNe, it is useful to keep in mind thefollowing isotopic characteristics compared to the solar isotope composition.Grains from supernova (type SN II) ejecta should show strong enrichments in the

stable isotopes12

C,15

N,28

Si, and excesses in26

Mg due to radioactive decay of 26Al, 41K from 41Ca, 44Ca from 44Ti, and 49Ti from 49V (Amari et al. 1996, Nittler et al. 1996, Travaglio et al. 1999, Hoppe et al. 2000, Hoppe and Besmehn 2002).Among these isotopes, only a few (i.e., 28Si, 44Ti, 49V) are unique products of and

52



tracers for SNe (including SN Ia for 28Si and 44Ti), whereas production of someof the other isotopes can also happen in other stellar sources (e.g., 26Al and 41Cain AGB and WR stars).

7.3.1. SiC from supernovae: X grains

Only 1% of all SiC grains are from supernovae (Amari et al. 1992, Nittler etal. 1996, Travaglio et al. 1999, Hoppe et al. 2000). The X grains have large 28Siexcesses (i.e., low 29Si/ 28Si and 30Si/ 28Si ratios), and SNe are the main source of 28Si. Further evidence for its supernova origin comes from grains that carry 44Caexcesses from decay of 44Ti (t 1/2= 60 a), which is only produced during explosivenucleosynthesis in SNe. Nittler et al. (1996) found that SiC X grains (as well assome low-density graphite grains) show evidence for 44Ti in the form of 44Caexcesses (up to 138 × solar) and that all grains (except one graphite grain) showcorrelated 28Si (up to 2 × solar) and 44Ca excesses. Besmehn and Hoppe (2003)found elevated 44Ca/ 40Ca ratios and 28Si excesses in ~20% of all SiC type Xgrains that they analyzed. The high 12C/13C ratios (>100) and relatively low14 N/ 15 N ratios (~20 to ~200) in X grains (Fig. 6) are consistent withnucleosynthesis in the He/C and He/N zones (Figs. 23,24), although not enough15 N is predicted in the SN models. The 26Al/ 27Al ratios in X grains of up to 0.6indicate contributions from the He/N zone to the supernova mixture from whichthe SiC grains condensed.

Trace element abundances in two X grains are very low (Fig. 9). Assumingrelative solar abundances for the elements condensing into the SiC X grains, theobserved abundance patterns can be explained if 98% of all trace elements wereremoved into other condensates such as refractory trace element carbides beforethe SiC condensed (Lodders and Fegley 1995). This seems plausible because

graphite grains ascribed to a supernova origin contain Ti-carbide and metalsubgrains (e.g., Croat et al. 2003). However, there is also the possibility thatelemental abundances in the gas from which SiC condensed were modified by r -and p-process nucleosynthesis. Future trace element analyses of supernova SiCgrains are required to sort out the different possibilities.

Isotopic ratios of Zr, Mo, and Ba in X grains have been analyzed by RIMS(Pellin et al. 2000a,b). Since the r -process takes place in SNe, one would expectthat isotopes of heavy elements in X grains show signatures of the r -process.However, the expected r -process signature was not found for Mo in X grains.Four out of 6 X grains have excesses in 95Mo and 97Mo whereas excesses in100Mo are expected from r -process models. Inspired by the new data andrevisiting predictions by Howard et al. (1992), Meyer et al. (2000) reconstructed

the neutron burst model. It postulates a rapid release of neutrons (on a time scaleof seconds) in He-rich matter heated by the passage of the shock wave whichgives a neutron flux lower than that of the classical r-process. This modelsuccessfully explained the 95Mo and 97Mo excesses of the grains. Zirconium and

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Ba isotopic ratios of the X grains can be also explained by the neutron burstmodel, assuming the grains condensed in a time scale of a few years.

7.3.2. Silicon nitride grains from supernovae

The few Si 3 N4 grains that have been discovered share many properties withthe SiC X grains and are therefore also related to supernovae (Nittler et al. 1995,Hoppe et al. 2000, Lin et al. 2002, Besmehn and Hoppe 2003)

7.3.3. Graphite grains from supernovae

Many of the low-density graphite grains are believed to have formed in SNe.They have similar isotopic compositions as SiC type X and Si 3 N4 grains: high12C/13C (>100) (Fig. 11), high 26Al/ 27Al up to 0.2 (Fig. 8), and Si-isotopicanomalies. Unlike SiC and Si 3 N4 of SN origin, graphite grains containmeasurable amounts of O and often have high 18O/16O ratios - up to 184×solar

(Fig. 14). The high26

Al/27

Al ratios (Fig. 8) are consistent with nucleosynthesis inthe He-N zone (Fig. 24), and the high 18O/16O ratios in the low-density graphites(Fig. 14) are a signature of nucleosynthesis in the He-C zone (Fig.24).

The observed isotopic composition of graphite grains requires mixing of different compositional SN zones because no single zone can explain theobservations. Travaglio et al. (1999) performed mixing calculations using thedifferent compositions of zones from Woosley and Weaver (1995) toquantitatively reproduce isotopic data for the low-density graphite grains. Thesemodels reproduced the observed 12C/13C, 18O/16O, and 30Si/ 28Si ratios and theinferred 41Ca/ 40Ca and 44Ti/ 48Ti ratios if jets of material from the inner Si-richzone penetrated the intermediate O-rich zones and mixed with material from theouter C-rich zones. However, the major problems are that the models do not

produce enough15

N and29

Si (e.g., Nittler et al. 1995, Travaglio et al. 1999,Hoppe et al. 2000).

7.3.4. Diamonds from supernovae

The essentially solar C- and N-isotopic compositions of meteoritic nanodiamondsare not helpful to relate them to any particular stellar source. On the other hand,the noble gas components with signatures of the r - and p- process leave littledoubt that SNe contributed to the presolar diamonds, which leaves the diamonds’C- and N-isotopic compositions somewhat of a puzzle. If all nanodiamonds werefrom SNe, it would mean that the highly variable C-isotopic compositions of thedifferent burning zones mix in such a manner as to give a 12C/13C ratio close tothe solar ratio. This would also have to happen in many SNe because likely morethan one contributed to the diamonds.

We noted before that Xe-HL in presolar diamonds is enriched in p- process124Xe and 126Xe and r -process 134Xe and 136Xe (Fig. 16). Hence the conclusionthat SNe are the places where the Xe-HL was implanted into the carriers. Themid-IR spectra of SN 1987A show a broad feature at 3.40 and 3.53 µ m (Meikle

54



et al. 1989), consistent with the identification of surface hydrogenated diamond(Guillois et al. 1999). This also supports the SN origin of some of the presolar diamond. Tielens et al. (1987) proposed that diamonds form by transformation of

preexisting graphite in supernova shock-waves. Huss and Lewis (1994b) andRussell et al. (1996) argue that this may not work because the diamonds shouldretain the noble gases that are already in the graphite grains; however, graphitegrains with supernova signatures do not contain Xe-HL, so diamonds could notinherit it. However, it may not be necessary for the graphites to already containthe Xe-HL component. Very massive stars go through the Wolf-Rayet (WR)stage, and in C-rich circumstellar shells of WR stars, graphite may condense. Allthis happens before the star goes supernova, and it is only during the supernovawhen the r -and p-processes responsible for Xe-HL are operating. Consequently,graphite formed earlier cannot contain Xe-HL because the r - and p-processes hadnot yet happened. However, once the products of the r -and p-processes ejectedwith the SN shock wave hit the older WR ejecta, the Xe-HL and other r -and p-

process products may be incorporated at the same time as the preexisting graphitegrains are transformed into diamonds.

Several questions remain about the details of the production of the p- and r - process nuclides of Xe seen in the nanodiamonds. Standard p- and r -processmodels produce equal excesses over solar for the p-process 124Xe and 126Xe, aswell as for r -process 134Xe and 136Xe, but in Xe-HL the excesses are not equal for each pair (Fig. 16). In order to explain the heavy Xe isotopes (“Xe-H”), Howardet al. (1992) proposed a neutron-burst model. This model assumes that a neutron

burst occurs in the He-rich zone of the supernova when this zone is heated (toabout 10 9K) by the passage of the shock wave (Meyer et al. 2000). This may leadto lower neutron-densities than in the classical r -process, so that 134Xe and 136Xeare produced with different efficiencies. However, the predicted 134Xe/ 136Xeratios from this model are different from those observed in Xe-H.

Remaining in the framework of the standard r -process, Ott (1996) suggestedthat a separation of the Xe isotopes and their radioactive precursors may occur

before decay of their precursors is completed. All precursors of 136Xe have half-lives on the order of one minute or less, while 134Te and 134I (which decay to134Xe) have half-lives of 42 and 52 minutes, respectively. In order to explain the134Xe/ 136Xe ratio, Ott (1996) estimated the timescale of the separation to be onthe order of hours. This scenario qualitatively accounts for the 124Xe/ 126Xe of Xe-L as well. The light isotopes 124Xe and 126Xe are produced from the precursors124Ba and 126Ba, with half-lives of 12 and 100 minutes, respectively. If theseparation between the Xe isotopes and their Ba precursors took place withinhours after the p-process, 124Xe would be more enriched than 126Xe. Analyses of Te isotopes in diamond (Richter et al. 1998, Maas et al. 2001) show that the ratioof the two r -process isotopes, 128Te/ 130Te, is also consistent with the “rapid”separation model proposed by Ott (1996). However, this requires a lot of fine

55



tuning for the timing of the parent-daughter nuclide separation during a veryturbulent SN mixing event. It remains to be seen if this can be a realistic model.

7.3.5. Oxide grains from supernovae

It is puzzling why presolar oxide grains of SN origin are so few if SNe areefficient dust producers. One possible explanation is that supernova oxide grainsare too small (<<0.1 µ m) to be recovered from meteorites during the presolar grain separation procedures. Since 16O is the third most abundant isotope ejectedfrom SNe, and the overall C/O ratio of the ejecta has C/O < 1 (Woosley andWeaver 1995), we would expect that oxides condense, and, for the same reason,that supernova oxides show larger excesses in 16O. However, there is only onecorundum grain (shown at the bottom of group III in Fig. 14) with a significant16O excess, which may have a supernova origin (Nittler et al. 1998). One spinelgrain, for which Choi et al. (1998) suggest a possible supernova origin, exhibitsan 18O excess (shown in Fig. 14 under the label for group IV). Otherwise, its O-isotopic composition is not very different from the presolar oxides of group IV,and its origin in a supernova appears uncertain. However, the 18O/16Ocomposition of graphite grains with SN signatures are quite similar to the 18O/16Oratios of group IV oxide grains of uncertain origin (Fig. 14). Measurements of other isotopes for this spinel do not provide better clues. The 26Al/ 27Al is only6×10 -4, far below the ratios (of up to ~0.1 26Al/ 27Al) seen in graphite and SiCgrains of clear supernova origin (Fig. 8). Except for an excess in the neutron-richisotope 50Ti, no isotopic anomalies in Ti and Ca were found by Choi et al. (1998),again different from what other presolar minerals from SNe show. On the other hand, such comparisons are somewhat complicated because it is not entirely clear to what extent the different zones, and hence the elements and isotopes in the

supernova are mixed, and, from which “supernova mixtures” the different presolar supernova grains are coming.

8. Binary stars as presolar dust sources

Novae and supernovae of type Ia are dust sources involving stellar binaries.These sources may not be unimportant because ~20% of stars occur in binary or higher multiple systems. Most stars, including those in binaries, are of lower mass (< 9 M ? ) and therefore evolve from the main sequence to giant stars andend up as white dwarfs. Of two stars, born at the same time but with somewhatdifferent masses, the more massive one evolves faster, and, depending on their relative age and the distance from each other, there are different possibilities for mass transfer between them.

In the case where one star is still a main sequence star while the companionstar is already a white dwarf, envelope mass from the main sequence star mayaccrete onto the hot white dwarf, which leads to a nova. A few presolar grains

56



seem to come from such systems. The transferred mass can be assumed to be of normal solar composition, and the white dwarf is either C and O-rich (a CO-WD)or O and Ne -rich (a ONe-WD) with masses ranging between 0.65 to 1.35 M ? .When the white dwarf has accreted enough matter to exceed the Chandrasekhar mass limit (1.4 M ? ) for thermonuclear explosions, a nova occurs.

The accreted material undergoes nuclear processing before it is ejected fromthe white dwarf. Dust condensation then incorporates the products of novanucleosynthesis. Models predict that isotopes of light elements up to Ca are

processed. In particular, low 12C/13C (<<10) and high 26Al/ 27Al ratios areexpected in ejecta of CO and ONe novae (Starrfield et al. 1997, 1998, José andHernanz 1998, José et al. 2001, 2004). In CO novae, the peak temperatures donot get high enough to significantly modify Si isotopic ratios, but enrichments in30Si are expected in novae of higher-mass ONe white dwarfs. A few presolar SiCand graphite grains with isotopic signatures consistent with theoretical

predictions for ONe novae are known (see Figs. 6-8 and 12, Amari et al. 2001c).Presolar grains from CO novae have not yet been identified, but the absence of graphite and SiC grains from CO novae is consistent with expectations fromcondensation calculations for nova ejecta (José et al. 2004).

Supernovae of type Ia (SN Ia) seem to result from interaction of two whitedwarfs. Explaining the thermonuclear explosions of SNe Ia by a merger of twowhite dwarfs in a close binary system can account for the absence of H in theejecta of SNe Ia and satisfies the requirement that SNe Ia should have low- tointermediate mass progenitors that do not evolve into supernovae as single stars(e.g., Iben and Tutukov 1984, Webbink 1984). Ejecta of SN Ia explosions shouldalso produce dust because many rock-forming elements are expelled. Clayton etal. (1997) suggested that the isotopic signatures of some SiC X grains areconsistent with an origin from SNe Ia caused by He accretion onto a CO-WD.However, it appears that the origin of these grains remains best understood by anorigin in mixed SN type II ejecta (Amari et al. 1998).

9. Conclusions and outlook

Much has been learned from presolar grains, but many questions remain and newones must be posed. Our collection of presolar minerals is not complete; theremust be other types of presolar minerals carrying abundant elements such as Fe,Cr, Mn, S, and P, just to name a few. It was about a year ago that the elusive

presolar silicates were first discovered, and the future may add to the array of

presolar minerals hidden in meteorites. Astronomical observations with a newgeneration of telescopes (e.g., the Spitzer telescope just taken into operation)already reveal amazing details of dust around giant stars and planetary nebulae.Also not long ago, observations at longer infrared and sub-millimeter wavelengths with ever-improving instruments provided the first evidence of the

57



previously “missing” dust from supernovae, and the detection of currently notobserved minerals around the different dust producing stars is awaiting. On thehand, advances in the developments of new micro- (or better nano-) analyticaltechniques such as the NanoSIMS and RIMS are directed towards obtainingmore detailed correlated measurements of mineralogical, chemical, and isotopic

properties of the known presolar components, which set the stage for probingnew components yet to be found. Such measurements, in turn, provide sensitivetests to stellar evolution and nucleosynthesis models. It will not be the last timethat the results from presolar grain measurements lead to revisions and updates insuch models, and prompt new measurements of physical properties, such asnuclear reaction cross-sections. Overall, the understanding of presolar grains andtheir implications requires combined efforts from astronomy, physics, chemistry,and mineralogy and will remain an exciting field of study in the future.

Acknowledgements : the authors thank Klaus Keil for the invitation for thisreview. They thank Roberto Gallino and Laura Schaefer for careful comments onthe manuscript. The authors also thank Gary Huss for comments. Work by K.L.was supported in parts by NASA grants NNG04GG13G and NNG04G157A fromthe NASA Astrobiology Institute. Work by S.A. was supported by NASA grants

NAG5-11545 and NNG04GG13G.

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